Patrick Moore’s Practical Astronomy Series

For further volumes:http://www.springer.com/series/3192

How to Observe the Sun Safely

Second Edition

Lee Macdonald

ISSN 1431-9756ISBN 978-1-4614-3824-3 ISBN 978-1-4614-3825-0 (eBook)DOI 10.1007/978-1-4614-3825-0Springer New York Heidelberg Dordrecht London

Library of Congress Control Number: 2012937044

© Springer Science+Business Media New York 2012This work is subject to copyright. All rights are reserved by the Publisher, whether the whole or part of the material is concerned, speci fi cally the rights of translation, reprinting, reuse of illustrations, recitation, broadcasting, reproduction on micro fi lms or in any other physical way, and transmission or information storage and retrieval, electronic adaptation, computer software, or by similar or dissimilar methodology now known or hereafter developed. Exempted from this legal reservation are brief excerpts in connection with reviews or scholarly analysis or material supplied speci fi cally for the purpose of being entered and executed on a computer system, for exclusive use by the purchaser of the work. Duplication of this publication or parts thereof is permitted only under the provisions of the Copyright Law of the Publisher’s location, in its current version, and permission for use must always be obtained from Springer. Permissions for use may be obtained through RightsLink at the Copyright Clearance Center. Violations are liable to prosecution under the respective Copyright Law.The use of general descriptive names, registered names, trademarks, service marks, etc. in this publication does not imply, even in the absence of a speci fi c statement, that such names are exempt from the relevant protective laws and regulations and therefore free for general use. While the advice and information in this book are believed to be true and accurate at the date of publication, neither the authors nor the editors nor the publisher can accept any legal responsibility for any errors or omissions that may be made. The publisher makes no warranty, express or implied, with respect to the material contained herein.

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Lee MacdonaldCambridge,UK

v

Preface

There is now more interest in observing the Sun than ever before among amateur astronomers. Go to any major astronomical meeting or convention and you will see at least one solar telescope in action, and it is sure to draw a crowd. The Sun indeed has much to offer the amateur astronomer with modest equipment. On most days, it shows sunspots and other features that display a wealth of fi ne detail and change their appearance strikingly from day to day. But observing the Sun can be danger-ous. Never look at the sun through an ordinary telescope or other optical aid, even for a brief instant . The Sun’s intense radiation, ampli fi ed and focused by a tele-scope, will almost certainly cause eye injury and could well lead to complete blind-ness. Do not attempt any solar observing until you have read and understood the safety precautions and observing advice set out in Chap. 2 of this book – even if you think you have the correct equipment. Be especially wary about using fi lters to observe the Sun. If you have a fi lter that makes the Sun look dark, it is not neces-sarily safe, as it is largely the Sun’s invisible radiation that is harmful to the eye. However, provided you use the correct techniques, such as projecting the solar image onto a screen or using a specially designed high-quality solar fi lter that fi ts over the telescope aperture, it is quite easy to observe the Sun safely.

One of the joys of solar observing is that useful observations are possible even with very small telescopes – such as the small refractors, Schmidt–Cassegrains, and Maksutov telescopes − that are readily available off the shelf. In fact, due in part to the fact that the Sun has more than enough light, a small telescope can actually give better results than a large one! Observing the Sun is also not affected by light pollu-tion, a major advantage for the many amateur astronomers whose view of the night sky is obscured by the glow of streetlights and security lighting. The Sun can be observed from a busy town just as successfully as from the remote countryside.

vi Preface

Our nearest star is studied intensively by professional astronomers and is moni-tored around the clock, using space-based observatories as well as telescopes from the ground, and both the level of research and the equipment required to carry it out are far beyond the amateur’s means. Therefore, solar observing does not offer the potential for discoveries or making major scienti fi c contributions, like some other branches of amateur astronomy, such as supernova hunting or variable star observ-ing. But monitoring solar features and keeping careful records of them is still important. Throughout the world, many amateur astronomers systematically moni-tor the Sun and send their observations to solar observing organizations for analy-sis. Monitoring levels of sunspot activity is particularly useful, as it continues a long series of observations made with small telescopes since the nineteenth century, which provide by far the best long-term record of solar activity that we have and is vital to our understanding of the Sun’s behavior and any effects it might have on Earth’s climate. More observers are always welcome in these sunspot counting programs. Solar photography is also useful, as it has considerable educational value. Professional solar images tend only to show small parts of the Sun or show our nearest star at invisible wavelengths, where its appearance is radically different from that in visible light. Amateur images, on the other hand, portray the Sun more realistically and so are more meaningful to the wider public. Indeed, as well as amateur astronomers, this book is also intended for those bringing astronomy to a wider audience, such as professional scientists engaged in public outreach activi-ties, which are increasingly important in the present age of budget cutbacks, when scientists are under increasing pressure to bring their subject to the public and jus-tify its value to the taxpayer.

The first edition of How to Observe the Sun Safely (Springer, 2003), was mostly written during 2001. But since that time, solar observing – and amateur astronomy as a whole – has undergone radical changes. The most fundamental of these has been the digital revolution and the almost complete substitution of digital imaging for fi lm photography. The first edition had a chapter on digital photography, but the book’s main emphasis was on 35-mm work. Digital cameras were still in their infancy: their resolution and exposure capabilities were modest, and they were dif fi cult to use with telescopes.

Digital SLRs did exist, but they cost over $2,000, putting them out of the reach of most amateurs. And no one had thought of using a webcam, costing (and weigh-ing) less than an eyepiece to take high-resolution images. All this has now changed. Digital SLRs now start at under $500 and take better images than their 35-mm predecessors, with all the bene fi ts of digital imaging − the ability to see and evalu-ate your results on the spot, and no more waiting to fi nish a roll of fi lm and have it developed.

At the same time, amateurs are routinely using webcams to take images of sun-spots and H-alpha solar features with a resolution once reserved for professional observatories. Another revolution has taken place in the affordability, and avail-ability, of telescopes and fi lters for observing the Sun in H-alpha. In 2001, the revo-lution was beginning, with the appearance of the Coronado SolarMax 40, the fi rst “sub-angstrom” H-alpha fi lter to be available for under $1,000. And, since then, the

viiPreface

revolution has continued. Now, some solar telescopes costing as little as $500 can show spectacular solar features that before 2001 would have required instruments costing fi ve to ten times as much. Two other companies – Lunt and Solarscope – have appeared and are producing high-quality H-alpha fi lters and telescopes, as has a revived DayStar (long the only source of sub-angstrom H-alpha fi lters for ama-teurs), with the result that the amateur now has a vast and potentially confusing array of equipment to choose from.

Therefore, there is all the more need for an up-to-date guide to show the amateur what to look for on the Sun, how to record observations, and what equipment to use. This second edition is aimed at the amateur who knows the basics of astronomy and wants to know how to go about observing the Sun. What is emphasized is what is possible using commercially available equipment that is easy to get hold of in most parts of the world. For this reason, I have deliberately eschewed some specialized topics, such as observing the Sun’s radio emissions, which requires homemade equipment and a fair amount of technical know-how. Neither do I discuss in much detail the Sun-related topics of eclipses and the aurora. Both are major fi elds in astronomy by themselves, and some good books on them have already been published.

Throughout the book, the emphasis is on practical solar observing – what you can do with ordinary equipment, provided you take the proper safety precautions. I have tried to avoid unnecessary theory and have not attempted detailed scienti fi c explanations, as these are available elsewhere. Rather, this book is intended as a basic guide to give the amateur a taste for observing our ever-changing nearest star, in the hope that he or she will explore further.

ix

Acknowledgments

I would like to thank Dave Tyler for permission to use his exceptionally detailed images – surely among the best solar images ever taken by an amateur astronomer – and Derek Hatch for his fi ne solar eclipse pictures. Thanks are also due to Richard Bailey (Solar Section Director of the Society for Popular Astronomy) for supplying photographs of his H-alpha fi lter.

I owe a particular debt of gratitude to Dr. Dominic Ford of the Cavendish Laboratory, Cambridge University, UK for his time in preparing the exquisite line diagrams from my freehand sketches, using the PyXPlot software that he wrote himself.

Another Cambridge colleague, Mark Hurn, kindly provided access to library facilities and enabled me to photograph some equipment at the Institute of Astronomy at Cambridge. Peter Meadows generously gave me permission to repro-duce one of his Stonyhurst solar disks, and Lyn Smith, Solar Section Director of the British Astronomical Association, allowed me to use BAA Solar Section data to plot a graph of solar activity from 2001 to 2010.

I am grateful to David Hathaway (NASA Marshall Space Flight Center) for permission to use the diagram in Fig. 4.8 . Last, but by no means least, I must thank John Watson, my editor at Springer for the first edition of How to Observe the Sun Safely , for starting off the process that has led to this new edition and Maury Solomon, Springer’s current Editor for Astronomy and Physics, for seeing this new book through the press.

Cambridge, UK Lee Macdonald June 2012

xi

Contents

1 Our Sun ...................................................................................................... 1The Sun’s Place in the Universe ................................................................. 2How the Sun Shines .................................................................................... 3The Sun’s Atmosphere ................................................................................ 4Solar Activity .............................................................................................. 6The Sun’s Influence on Earth ...................................................................... 12

2 Equipment for Observing the Sun ........................................................... 17The Sun’s Radiation .................................................................................... 17Telescopes for Solar Observing .................................................................. 19

The Refractor .......................................................................................... 19The Reflector .......................................................................................... 23Catadioptric Telescopes .......................................................................... 24

Telescope Mountings .................................................................................. 25Viewing the Sun’s Image ............................................................................ 27

Solar Projection ...................................................................................... 27Solar Filters ............................................................................................ 31Other Observing Methods ...................................................................... 36Observing the Sun with the Naked Eye and Binoculars......................... 37

3 What Can We See on the Sun? ................................................................ 41When to Observe the Sun ........................................................................... 41Where to Observe the Sun .......................................................................... 43Aiming the Telescope ................................................................................. 45Viewing the Sun’s Surface .......................................................................... 47

Granulation ............................................................................................. 47

xii Contents

Limb Darkening ..................................................................................... 48Sunspots .................................................................................................. 49Flares ...................................................................................................... 61

4 Solar Position Measurements ................................................................... 63Drawing Using the Projection Method ....................................................... 64

Making a Projection Grid ....................................................................... 64Orienting the Image ................................................................................ 67Making the Drawing ............................................................................... 68Deriving Sunspot Positions .................................................................... 71Example .................................................................................................. 74What We Can Learn from Drawings ...................................................... 78Detailed Drawings .................................................................................. 80

Cooperation with Other Observers ............................................................. 81

5 Measuring Solar Activity ......................................................................... 83The Mean Daily Frequency......................................................................... 84The Relative Sunspot Number .................................................................... 86Observing Faculae and White-Light Flares ................................................ 97Observing Naked-Eye Sunspots ................................................................. 98

6 Observing the Chromosphere .................................................................. 101Equipment for Observing the Chromosphere ............................................. 106

H-Alpha Telescopes................................................................................ 110H-Alpha Filters ....................................................................................... 114Choosing an H-Alpha System ................................................................ 121Calcium-K and Other Systems ............................................................... 124

Prominences and Filaments ........................................................................ 125Counting Prominences ........................................................................... 128Prominence Position Measurements ...................................................... 128

Flares ........................................................................................................... 131

7 Imaging the Sun with a Digital Camera ................................................. 135Choosing a Digital Camera ......................................................................... 137

“Compact” Cameras ............................................................................... 137Digital SLRs ........................................................................................... 142

Telescopes and Mounts ............................................................................... 144Filters .......................................................................................................... 146Mounting the Camera ................................................................................. 148Photographic Techniques ............................................................................ 150

Shooting the Projected Image ................................................................. 150The Afocal Method ................................................................................ 151Prime Focus Photography....................................................................... 151Using a Teleconverter ............................................................................. 154Eyepiece Projection ................................................................................ 156

xiiiContents

Taking Pictures ............................................................................................ 159Taking Pictures with a Compact Digital Camera ................................... 162Taking Pictures with a DSLR ................................................................. 164

Photographing the Chromosphere ............................................................... 168

8 Webcam Imaging and Image Processing ............................................... 173Webcams and Accessories .......................................................................... 174

Do You Need a Webcam? ....................................................................... 174Advantages of Webcams ........................................................................ 174Choosing a Webcam for Solar Imaging .................................................. 176Computers and Accessories .................................................................... 178

Taking Webcam Images .............................................................................. 179Processing Webcam Images ........................................................................ 182Enhancing Digital Images ........................................................................... 186

File Formats ............................................................................................ 188Cropping and Trimming ......................................................................... 190Image Orientation ................................................................................... 190Changing the Brightness and Contrast ................................................... 191Removing Dust and Scratches ................................................................ 192Sharpening the Image ............................................................................. 192Changing the Color ................................................................................ 193Making Composite Images ..................................................................... 195

Appendix A ................................................................................................. 197

Appendix B ................................................................................................. 203

Appendix C ................................................................................................. 205

Appendix D ................................................................................................. 207

Index ................................................................................................................. 211

xv

List of Figures

Fig. 1.1 Cutaway diagram showing the structure of the Sun ....................... 4 Fig. 1.2 Total solar eclipse, photographed by Derek Hatch

on August 1, 2008, showing the extensive white corona around the black silhouette of the Moon ............................ 5

Fig. 1.3 The Sun’s photosphere (“surface”), showing a very large sunspot group and several smaller groups, photographed by the author in August 2002 .................................. 7

Fig. 1.4 The planet Mercury in transit across the Sun, photographed by the author on May 7, 2003, with an 80 mm refractor and full-aperture solar filter. Note how the silhouette of Mercury, visible near the bottom of the picture, appears jet black, while the small sunspot near the top has a brownish tint, because sunspots appear dark only by contrast .............................. 8

Fig. 1.5 Total solar eclipse, August 1, 2008, showing prominences around the silhouette of the Moon. This is a short exposure of the same eclipse as in Fig. 1.2; it allows the prominences to show up without being drowned out by the bright inner corona (Photograph by Derek Hatch) ............................................. 10

Fig. 1.6 The Sun’s corona, imaged in the extreme ultraviolet by the SOHO satellite in September 1997, showing bright active regions and dark coronal holes. The active regions broadly correspond to sunspot groups in visible light. The photosphere appears dark because it is too cool to emit at this wavelength (Image courtesy of NASA) ............................................................. 11

xvi List of Figures

Fig. 1.7 A coronal mass ejection (CME) imaged in visible light by the SOHO satellite in November 1997. The Sun’s brilliant disc has been artificially hidden by an occulting disc inside the telescope (Image courtesy of NASA) ....................................... 12

Fig. 1.8 A major display of the aurora borealis, photographed from southern England by the author on April 6, 2000. The display was triggered by a major geomagnetic storm ............. 14

Fig. 2.1 The electromagnetic spectrum ........................................................ 18 Fig. 2.2 The author’s 80 mm (3.1 in.) refracting telescope on a German

equatorial mount with electronic slow motion controls. A home-made balsa wood projection box is attached to the eyepiece end of the telescope for safe viewing of the Sun’s image .......................................................................... 20

Fig. 2.3 The author’s 60 mm (2.4 in.) f/5.9 Takahashi apochromatic refractor. Its short tube makes it extremely compact and portable .................................................................................... 21

Fig. 2.4 A Meade ETX 90 mm (3.5 in.) Maksutov telescope equipped with a glass aperture filter. Note that the finderscope has been removed. This is an important safety precaution, because the Sun is just as dangerous to look at through the finderscope as at the main telescope, and careless people, particularly children, could look through it by accident .................................... 25

Fig. 2.5 Projecting the Sun’s image onto paper for safe solar viewing. A card has been fitted over the eyepiece end of the telescope to shade the image from direct sunlight ......................................... 27

Fig. 2.6 A solar projection screen, designed by Roderick Willstrop of Cambridge University, attached to the 200 mm (8-in.) Thorrowgood refractor at the Institute of Astronomy in Cambridge, England. Here the screen is made from a light framework of plywood – heavier than balsa wood – attached to a large and heavy telescope ........................................................ 29

Fig. 2.7 A balsa wood projection box made by the author and attached to an 80 mm (3.1-in.) refractor ....................................................... 29

Fig. 2.8 Aperture filter made from Baader AstroSolar Safety Film in a home-made mount, fitted to an 80 mm refractor ..................... 33

Fig. 2.9 A glass aperture filter mounted to fit the Meade ETX 90 mm (3.5 in.) Maksutov telescope pictured in Fig. 2.4 ........................... 35

Fig. 2.10 A Herschel wedge or Sun diagonal made by Lunt Solar Systems ................................................................................. 36

Fig. 2.11 A naked-eye solar viewer made from black polymer (left) and a pair of Mylar eclipse glasses (right) ..................................... 38

Fig. 3.1 Diagram showing the principle of a solar finderscope (Courtesy of Dominic Ford) ........................................................... 46

xviiList of Figures

Fig. 3.2 Solar finderscope on the author’s 80 mm refractor. This replaces the “normal” finderscope, which has been removed as a safety precaution ....................................................... 46

Fig. 3.3 Whole-disc photograph of the Sun, showing sunspots and limb darkening. The latter has been exaggerated by the high contrast of the image (Photograph by the author) ....... 48

Fig. 3.4 A small, simple sunspot, showing the dark central umbra and lighter surrounding penumbra (Image by Dave Tyler) ............ 51

Fig. 3.5 In this remarkably high-resolution image by Dave Tyler of a bipolar sunspot group, taken on April 23, 2011, the penumbral filaments are clearly visible in two of the group’s spots, especially the large spot on the right. The granulation of the surrounding photosphere also shows up strikingly .............. 51

Fig. 3.6 A small, symmetrical sunspot showing the Wilson Effect: the penumbra on the side of the spot nearest the center of the disc is narrower than that on the side nearest the limb, making the spot look like a depression on the Sun’s surface (Image by Dave Tyler) ................................................................ 52

Fig. 3.7 A large sunspot group, photographed by the author in July 2004, showing multiple light bridges in the umbra of its leader spot ............................................................................. 54

Fig. 3.8 Some examples of sunspot group types, ranging from simple to complex, showing their McIntosh classifications. All images by Dave Tyler. (a) Very small group of class Axx, (b) Small bipolar group, class Bro, (c) Spot group of class Dao, (d) Spot group of class Ekc, (e) Symmetrical spot, class Hhx ..................................................... 57

Fig. 3.9 The giant sunspot group of March 2001, McIntosh classification Fkc. This was one of the largest sunspot groups ever recorded and was associated with widespread auroral displays (Photograph by the author) .................................. 58

Fig. 3.10 Another of the giant sunspot groups of Solar Cycle 23, this time the largest of the groups of October 2003, McIntosh classification Fkc (Photograph by the author) ................ 58

Fig. 3.11 A sunspot group near the limb, showing extensive faculae, imaged by Dave Tyler on March 7, 2011. An image of the same group in H-alpha light, shown for comparison, appears at top .................................................................................. 60

Fig. 4.1 A solar drawing grid, consisting of a 152 mm (6 in.) circle divided into ½ in. squares, as used in the projection box and underneath the drawing. The 5° intervals marked out around the circumference are for determining the Sun’s true orientation for sunspot counting purposes (see Chap. 5) (Courtesy of Dominic Ford) ........................................................... 65

xviii List of Figures

Fig. 4.2 Close-up of the author’s solar projection box, showing the rotatable projection grid ............................................................. 67

Fig. 4.3 Orientation of the Sun’s image at midday as seen in various telescope arrangements. (a) Straight projection through a refractor or Newtonian (northern hemisphere). (b) Projection using a star diagonal (northern hemisphere). (c) Filtered solar image as seen directly in a refractor, Schmidt-Cassegrain or Maksutov telescope, using a star diagonal or 90° mirror (northern hemisphere). (d) Straight projection through a refractor or Newtonian (southern hemisphere). (e) Projection using a star diagonal (southern hemisphere). (f) Filtered solar image as seen directly in a refractor, Schmidt-Cassegrain or Maksutov telescope, using a star diagonal. In all arrangements with a star diagonal, these diagrams assume that the observer is standing or sitting directly behind the eyepiece and facing the front of the telescope tube (Courtesy of Dominic Ford) ................................................................................ 69

Fig. 4.4 An example of a Stonyhurst disc, for a solar tilt (B0) of 6°

(Courtesy of Peter Meadows) .............................................................. 72 Fig. 4.5 Illustrating the Sun’s changing position angle (P) and “nodding”

towards and away from us (B0) during the course of the year

(Courtesy of Dominic Ford) ........................................................... 73 Fig. 4.6 Whole-disc drawing, made by the author on October 12, 2008,

showing a small bipolar sunspot group. The position angle P is marked ..................................................................................... 75

Fig. 4.7 The author’s drawing of October 12, 2008, shown laid over a Stonyhurst disc representing B

0 = 6.0° (Stonyhurst disc courtesy

of Peter Meadows) .......................................................................... 77 Fig. 4.8 The “Butterfly Diagram,” plotted from the 1880s up to 2011

(top), matched up with the graph of sunspot numbers for the same time period (below). The migration of sunspots from high to low latitudes in each cycle is clearly seen (Courtesy of NASA/David Hathaway) ........................................... 79

Fig. 5.1 Whole-disc image of the Sun, showing three sunspot groups. The active area count for that day should therefore be recorded as 3 (Photograph by the author) ..................................................... 85

Fig. 5.2 This image of a sunspot group gives a good demonstration of which internal spots should and should not be counted as spots for the purpose of determining the Relative Sunspot Number, R. All the dark umbrae towards the left and right ends of the group should be counted, but the pores and gray penumbral material at the center should not (Image by Dave Tyler) ............... 89

Fig. 5.3 Diagram showing how to determine the Sun’s true orientation by eye. Note also that the Sun’s true equator is usually curved northwards or southwards, and this needs to be taken into

xixList of Figures

account when determining which hemisphere a spot is in (Drawing by Dominic Ford) ......................................... 91

Fig. 5.4 Sample page from the author’s solar notebook, showing sunspot counts recorded in tabular form ........................................ 92

Fig. 5.5 Graph plotted from the author’s sunspot counts, showing the variation in the MDF between 2001 and 2010 ......................... 95

Fig. 5.6 Graph plotted from the author’s sunspot counts, showing the variation in the Relative Sunspot Number R between 2001 and 2010 ................................................................................ 95

Fig. 5.7 Graph showing the variation in the Relative Sunspot Number R between 2001 and 2010, based on British Astronomical Association collective sunspot data (Permission to use BAA data kindly supplied by Lyn Smith, BAA Solar Section Director) ................... 96

Fig. 6.1 Diagram of the visible light spectrum, showing the wavelengths of some major chromospheric emission lines (Diagram by Dominic Ford) ................................................................................ 102

Fig. 6.2 Example of how the level of white-light sunspot activity is no accurate guide to what might be visible in H-alpha. This giant solar prominence was photographed by the author on April 4, 2004, when sunspot activity was very moderate (see text). Picture taken using a Canon EOS 300D digital SLR camera attached to an 80 mm refractor with Baader “H-alpha coronagraph” prominence viewer. Exposure 1/250 s with the camera set to ISO 400 .................................................................... 103

Fig. 6.3 The Sun’s disc in hydrogen-alpha. The dark linear features are filaments (i.e., prominences seen in absorption against the disc), and the bright regions are plages. Whole-disc mosaic image by Dave Tyler ................................................................................. 104

Fig. 6.4 An example of a “filaprom” – a filament seen in absorption against the solar disc extending beyond the Sun’s limb and becoming a prominence seen in emission against the apparent blackness of space (In reality, it is seen against the corona, but the latter is much too faint to be seen in an H-alpha filter.) (Image taken by Dave Tyler on November 20, 2009) .................... 105

Fig. 6.5 The Coronado Personal Solar Telescope (PST) – the 40 mm dedicated H-alpha telescope that made sub-angstrom H-alpha views of the Sun available for under $500. The large dark sheet is a Sun shade to shield the observer from direct sunlight (Photograph by the author) ................................................................ 110

Fig. 6.6 The Lunt Solar Systems 152 mm (6 in.) dedicated H-alpha solar telescope. In the background are Lunt’s 50 and 100 mm models (Photograph by the author) ............................................................. 112

Fig. 6.7 The Solarview 50 H-alpha telescope made by Solarscope Ltd. (Photograph by the author) ...................................................................... 113

xx List of Figures

Fig. 6.8 A 40 mm Coronado SolarMax filter fitted to the author’s Takahashi FS-60C 60 mm apochromatic refractor. (Photograph by the author) ............................................................. 114

Fig. 6.9 A DayStar T-Scanner sub-angstrom H-alpha filter, a popular model with amateur astronomers for many years and still used by many solar observers. (a) The filter unit housing the etalon. (b) The filter unit mounted on a Meade SCT (Photographs by Richard Bailey) .......................................................................... 118

Fig. 6.10 Attaching a Moon filter in front of the eyepiece will increase the contrast of disc features in an H-alpha image – though it does not narrow the passband and the prominences will be fainter (Photograph by the author) ............................................................. 119

Fig. 6.11 A Baader H-alpha coronagraph on an 80 mm (3.1 in.) refractor (Photograph by the author) ............................................................. 120

Fig. 6.12 A typical hedgerow prominence (Image by Dave Tyler) ............... 126Fig. 6.13 H-alpha image of a sunspot group, showing an active region

filament. Note that this type of filament is smaller and darker than typical quiescent filaments (Image by Dave Tyler) ................ 127

Fig. 6.14 H-alpha image of a flare in the large sunspot group AR11158, taken by Dave Tyler on February 17, 2011. The flare is the intensely bright region near the center of the group ....................... 132

Fig. 6.15 A spectacular limb flare (Imaged by Dave Tyler on March 8, 2011) .......................................................................... 132

Fig. 7.1 A typical compact digital camera: a 10-megapixel Canon PowerShot A640. This camera’s LCD viewfinder screen unfolds from the main body of the camera and can be adjusted to a great variety of angles – a handy feature when the camera is mounted on a telescope and the screen is at an awkward angle if left in its conventional position ...................................................................... 138

Fig. 7.2 A heavily compressed JPEG image, occupying only about 100 KB on the computer’s hard drive but rendered almost useless by the variation in tone caused by the limb darkening being reduced to a series of layers. The original image was taken on a 6.3-megapixel DSLR, demonstrating that the degree of image compression is more important than the size of the image in megapixels ...................................................................................... 140

Fig. 7.3 The author’s 6.3-megapixel Canon EOS 300D digital single lens reflex (DSLR) camera, shown with 18–55 mm zoom lens separated from the camera body. With the lens removed, cameras of this type can be securely clamped to the telescope with suitable adapters to take high-quality images of the Sun (and other astronomical objects) ...................................................................... 143

Fig. 7.4 A filter made from Baader AstroSolar Photo Film, in a home-made mount and fitted to the author’s 80 mm refractor.

xxiList of Figures

This type of filter transmits a brighter image than visual filters and so allows very short exposures to be used with DSLR cameras. All photographic filters should carry a warning label like this one, indicating that they are not safe for visual use .......... 147

Fig. 7.5 Three methods of using a compact digital camera at the telescope eyepiece to image the Sun. (a) A Canon PowerShot A640 camera held up to the eyepiece by hand. (b) The PowerShot A640 mounted on a tripod and pointed into the eyepiece. (c) A Nikon Coolpix 900 camera clamped to the eyepiece using a commercial digital camera adapter ................................................. 148

Fig. 7.6 Prime focus photography with a DSLR: the camera body is attached directly to the telescope drawtube with an adapter ...... 152

Fig. 7.7 Prime focus photography can be a good method for imaging a partial solar eclipse. This image of the partial eclipse seen just after sunrise from the UK on January 4, 2011, was taken by the author using a Canon 300D at the prime focus of a Vixen 80 mm refractor (focal length 910 mm), with a Baader AstroSolar Photo Film (ND3.8) filter. Exposure 1/125 s at ISO 400 ............... 153

Fig. 7.8 A small telescope of short focal length gives only a very small solar image when the camera is used at the prime focus. Image taken by the author on March 14, 2010, using a Canon 300D DSLR at the prime focus of a Takahashi FS-60C 60 mm refractor, focal length 355 mm, equipped with a full-aperture Baader AstroSolar (ND5) filter. Exposure 1/2,000 s .................................. 155

Fig. 7.9 Whole-disc solar image, showing a large sunspot group. Image taken by the author on September 13, 2005, using a Canon 300D DSLR on an 80 mm f/11.4 refractor with a 1.4× teleconverter, giving an effective focal length of 1,274 mm. The Sun’s disc nicely fills the frame in many DSLRs using this sort of focal length. Filters used were Baader full-aperture AstroSolar Photo Film (ND3.8) and No. 8 (light yellow) secondary filter. Exposure was 1/3,200 of a second at ISO 100 .................................................................... 156

Fig. 7.10 Diagram showing the principle of eyepiece projection using a refractor or Newtonian telescope, using a camera adapter and T-ring ....................................................................................... 157

Fig. 7.11 Eyepiece projection arrangement on the author’s 80 mm refractor, showing Canon EOS 300D DSLR (with cable release), T-ring and camera adapter ................................................................................ 157

Fig. 7.12 Close-up of two sunspot groups, an enlargement of an image photographed by the author on December 4, 2005, using an 80 mm refractor and eyepiece projection with a 15 mm eyepiece to give an effective focal length of 5,763 mm at f/72. Filters used were Baader full-aperture AstroSolar Photo Film (ND3.8) and No. 8 (light yellow) secondary filter. Exposure was 1/1,250 s at ISO 400 on a Canon 300D DSLR ................................................................. 159

xxii List of Figures

Fig. 7.13 The Canon “Angle Finder C” focusing magnifier fitted to a Canon 300D camera mounted on the author’s 80 mm refractor .......................................................................................... 166

Fig. 7.14 Partial solar eclipse, March 29, 2006, focused using the method for focusing a DSLR described in the text. Image taken with Canon 300D on 80 mm refractor with full-aperture Baader AstroSolar Photo Film (ND3.8) filter and 1.4× teleconverter. Exposure 1/3,200 of a second at ISO 100 ...................................... 167

Fig. 7.15 The Sun in H-alpha, taken by the author with a Canon PowerShot A640 hand-held to a 25 mm eyepiece on an 80 mm refractor equipped with a Coronado SolarMax 40 filter. This image demonstrates the difficulty in obtaining good H-alpha images with color cameras. Two filaments, some plage detail and a sunspot are faintly visible ............................................................... 170

Fig. 7.16 Prominences photographed through a coronagraph attached to an 80 mm refractor. A finely detailed prominence is visible to the upper right of this image originally shot on 35 mm film ...... 171

Fig. 8.1 The author’s 80 mm refractor with webcam connected to laptop computer .............................................................................. 175

Fig. 8.2 Celestron NexImage astronomical webcam attached to the author’s 80 mm refractor ...................................................... 176

Fig. 8.3 The tiny Flea3 webcam-type CCD camera on Dave Tyler’s 130 mm Astro Physics refractor mounted “piggyback” on his 14-in. Celestron SCT. The Flea3 is smaller than the Barlow lens it is attached to! ....................................................................... 178

Fig. 8.4 A single frame from a webcam movie has a grainy appearance and low resolution. This image of a group of prominences was taken by the author on February 11, 2008. It is one frame from a 900-frame movie taken with a Celestron NexImage webcam attached to a 60 mm refractor with a Coronado SolarMax 40 H-alpha filter ............................................................. 183

Fig. 8.5 The same prominences as in Fig. 8.4, after the webcam movie has been through the stacking and aligning process ....................... 185

Fig. 8.6 The final image of the prominences, after wavelet processing ....... 186 Fig. 8.7 Steps in enhancing a digital image of the Sun using Adobe

Photoshop Elements. (a) Original image, taken by the author on March 8, 2011, with a Canon 300D on an 80 mm refractor with a 1.4× teleconverter and a Baader AstroSolar Photo Film filter, (b) The image cropped to center the Sun’s image and exclude some of the black background, (c) The image flipped to show east to the right, (d) The brightness decreased by 20 points and contrast increased by 40 points, (e) Image after applying “Sharpen” twice and then with contrast and sharpness enhanced further using “Unsharp Mask,” (f) Image with red and green channels adjusted to give a distinct yellow color ........................................... 189

xxiiiList of Figures

Fig. 8.8 Whole-disc solar image, originally taken by the author on June 19, 2001, on Kodak Technical Pan 2415 black-and-white film, then more recently scanned and tinted yellow using Adobe Photoshop Elements ............................................................ 195

Fig. 8.9 The webcam image of the Sun shown in Fig. 8.6, with the overexposed solar disc blacked out using Adobe Photoshop Elements, mimicking the effect of a coronagraph or a total solar eclipse ............................................................................................. 196

Fig. A.1 Diagram showing the construction of the author’s projection box. The projection distance d determines the diameter of the projected solar image ...................................................................... 199

Fig. A.2 The author’s projection box, attached to an 80 mm (3.1 in.) refractor with a pipe bracket ........................................................... 201

xxv

About the Author

A passionate amateur astronomer since childhood, Lee Macdonald has been observing the Sun for more than a quarter of a century and has witnessed the rise and fall of two solar cycles. Since the late 1980s, he has reported his solar observa-tions to the British Astronomical Association and the Society for Popular Astronomy (originally the Junior Astronomical Society). Since the late 1990s, he has also observed the Sun in H-alpha.

Lee is also an experienced writer on astronomy. As well as the fi rst edition of How to Observe the Sun Safely , Lee has written articles for numerous magazines, including Astronomy and Astronomy Now . He is an experienced deep-sky observer and for four years served as Editor of The Deep-Sky Observer , the magazine of the Webb Deep-Sky Society.

A historian by training, Lee has degrees in History and Modern History from the University of Reading in the UK and a third degree in History and Philosophy of Science from the University of Cambridge. He specializes in the history of astron-omy and has published two major research papers on the history of the 2.5-m Isaac Newton Telescope – one of the world’s best-known telescopes, now on La Palma in the Canary Islands. He has also published a paper on astrophotography pioneers Isaac Roberts and Edward Emerson Barnard in the prestigious Journal for the History of Astronomy . In 2008, he was elected a Fellow of the Royal Astronomical Society.

Lee currently works as an administrator and secretary in the Department of Applied Mathematics and Theoretical Physics at the University of Cambridge, working for Professor Stephen Hawking’s relativity and cosmology research group.

1L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0_1, © Springer Science+Business Media New York 2012

Our Sun

Chapter 1

The Sun is important to astronomers for two reasons. The fi rst is that it is Earth’s only natural “power station,” producing the light and heat essential to life on our planet. Without the Sun, Earth would be more or less a frozen ball of rock, with no atmosphere, no weather, no life and no people. The Sun can also be harmful to us, because it emits huge quantities of radiation that would be fatal to humans and all living matter, were our planet not protected from it by a thick atmosphere and pow-erful magnetic fi eld. But intense bursts of solar activity can still harm communica-tions and electrical power systems, so we need to keep a constant watch on the Sun so that we are forewarned of its next powerful outburst. We also need to understand it so that we can predict future activity and its likely consequences.

The second reason for studying the Sun is that the Sun is a star, much like the 3,000 or so other stars that we can see in the night sky. All the stars in the night sky, however, are exceedingly remote from us. Even the closest known star, Proxima in the southern constellation of Centaurus, is some 60 million million km from Earth – so far away that its light takes over 4 years to reach us. We can say, therefore, that Proxima Centauri is over 4 light-years from us, a light year being the distance trav-eled by light in 1 year – approximately 9.46 × 10 12 km. Most stars – even the major-ity of those visible to the naked eye – are much further away still. Even the world’s most powerful telescopes show the stars as mere points of light and give us just basic information as to what they are and how they work. But at only 150 million km from Earth, the Sun is easily close enough for us to have a detailed view and learn much about it. Armed with a detailed knowledge of the Sun, astronomers can learn more about other stars.

2 1 Our Sun

The Sun’s Place in the Universe

The Sun in the sky may seem quite small – about half a degree in angular diameter, the same size as the Moon. But in human terms, the Sun is vast: 1.4 million km in diameter and with a volume large enough to contain over a million Earths. It is, of course, by far the brightest object in our sky, some 400,000 times brighter than the full Moon and emitting an incredible 4 × 10 26 W of radiation.

However, compared to other stars in the universe, the Sun is an average star in terms of its size and luminosity. Some stars are much larger and brighter than the Sun and some of the largest, the “supergiants,” would extend out to the orbit of Mars or even beyond if placed in the Solar System. At the other end of the scale, the “dwarf” stars emit a relatively feeble light and are typically no larger than Earth. The varying amounts of brightness of the stars that we see in the night sky give no indication of the stars’ sizes and actual brightness, because the apparent brightness of a star is also determined by its distance from us. Some stars in the night sky appear bright only because they are relatively close to us. A good example is Sirius, the brightest star in the sky. Sirius is, in fact, quite an average star, but it is quite close by at “only” 9 light-years. However, not far away in the winter sky is Rigel, at the foot of the constellation Orion. It appears only slightly fainter than Sirius, but it is 800 light-years away, making it 60,000 times as lumi-nous as the Sun!

You can get a good analogy of the difference between a star’s true and apparent brightness by standing in the same room as a 40-W light bulb and looking out at a distant fl oodlight or security light. The light bulb a few feet away from you appears brighter, even though the security light is a much more powerful device. As stars go, the Sun is one of the “light bulbs” – not one of the dimmest stars in our galaxy but certainly not one of the brightest, either.

In ancient times – and, indeed, until less than 500 years ago when Copernicus published his “heliocentric” theory of the universe – it was generally believed that Earth lay at the center of the universe, and all the other heavenly bodies – including the Sun – circled around it. Astronomical research in the fi ve centuries since Copernicus has revealed that nothing could be further from the truth. Our Earth and even the Sun are very junior members of an exceedingly complex universe. Earth is but one of eight planets (plus countless smaller bodies) orbiting the Sun, a very average star. The Sun is just one of 100 billion (in this book we shall take ‘billion’ to mean a thousand million, or 10 9 ) stars in a huge, rotating disc that we call the Milky Way Galaxy. The Sun and Solar System are nowhere near the center of the galaxy; in fact, we are about 28,000 light-years out, in one of the spiral arms. The Milky Way, in turn, is just one of millions of galaxies scattered throughout the universe. The Sun is thus relatively unimportant in the grander scheme of things, but to us on Earth it is essential to our existence.

3How the Sun Shines

How the Sun Shines

The Sun is a gaseous sphere made mostly of hydrogen and helium. By mass, 71% of the Sun is hydrogen and 28% helium, with traces of many other elements making up the remaining 1%. The Sun and other stars are nuclear powerhouses that gener-ate their own heat energy by means of nuclear fusion reactions at their cores. Nuclear fusion is a reaction in which atomic nuclei of one element combine together to form another, “heavier” element – that is, one with a larger number of protons and neutrons and further down the Periodic Table of the Elements. At the core of the Sun, protons – which are simply the nuclei of hydrogen atoms, a hydro-gen atom consisting of just one proton and one electron – move around at extremely high speeds, and occasionally two exceptionally fast-moving protons will collide and merge together. In a sequence of such collisions, known as a “proton-proton chain,” four protons fuse together to create a single nucleus of helium. In this fusion reaction, 0.7% of the total mass of the four hydrogen nuclei is lost, because it is converted to energy in accordance with Albert Einstein’s most famous equation:

2E mc=

where E is the energy produced, m is the mass lost in the reaction and c is the speed of light. Because c is vast (approximately 300,000 km per second), the amount of energy released is huge, and it is this energy that powers the Sun and stars. Through this fusion process, every second the Sun converts 5 million metric tons of hydro-gen into energy, which is eventually released into space. The Sun is thus continually getting lighter, but the mass loss is so small in relation to the Sun’s total mass that it will have no signi fi cant effect for billions of years to come. For fusion to happen, the temperature at the core has to be very high. Current estimates say that the Sun’s core has a temperature of around 16 million degrees Kelvin (commonly written as 16 million K). 1

The core in which the fusion reactions occur extends outwards to about a quarter of the Sun’s radius (Fig. 1.1 ). The energy produced by the core is radiated outwards to the upper regions of the Sun’s interior through a very stable region of gas extend-ing to more than two-thirds of the way to the Sun’s surface. In this radiative zone energy from the core is constantly absorbed, re-radiated and de fl ected by the hot surrounding gas, so that it takes some 170,000 years for energy to travel from the core to the upper regions of the Sun. Thus the energy that we see as light on the Sun’s surface today started its journey from the core during the last Ice Age!

1 Kelvin is the temperature measured from “absolute zero,” approximately −273°C, and is the temperature scale most commonly used by astronomers. To convert from Kelvin to Celsius, sub-tract 273°. At the huge temperatures we are dealing with in the Sun, the difference between Kelvin and Celsius is insigni fi cant.

4 1 Our Sun

At about 500,000 km from the Sun’s center the gas becomes too cool for the radia-tion to fl ow outwards via the above process. As the radiation is blocked, pressure builds up, and this causes the radiation to fl ow to the surface via convection currents, much as heat from a radiator or fi re rises into the cooler surrounding air. Heat takes about 10 days to pass through this convective zone , at the end of which it fi nally arrives at the surface, known to astronomers as the photosphere , which we see as the luminous disc of the Sun. (“Photosphere” means “sphere of light.”) Although the photosphere appears brilliant to the eye, it is actually relatively cool, at “only” 5,800 K.

Actually, the word “surface” is a misnomer here, since it is not solid but gaseous, as is the rest of the Sun. However, the photosphere is a surface in the sense that it sharply de fi nes the outline of the Sun as seen in ordinary, visible light. It is also the layer from which most of the Sun’s energy is radiated.

The Sun’s Atmosphere

The Sun is surrounded by a very extensive atmosphere, but this atmosphere is extremely thin and much fainter than the photosphere. As seen from Earth, the Sun’s atmosphere is totally overwhelmed, because the brilliant light of the photo-sphere is scattered by Earth’s atmosphere. The solar atmosphere is therefore invis-ible from Earth except when the photosphere is hidden by the Moon during a total solar eclipse (Fig. 1.2 ). The Sun is some 400 times the diameter of the Moon, but it is also 400 times further away from Earth, so it subtends the same apparent size as seen from our planet. At new Moon – that is, when the Moon passes between Earth and the Sun in its monthly orbit – the Moon very occasionally passes right in

Convective Zone

Radiative Zone

Core

PhotosphereSunspots

Fig. 1.1 Cutaway diagram showing the structure of the Sun

5The Sun’s Atmosphere

front of the Sun, causing an eclipse. Solar eclipses only happen once or twice a year anywhere on Earth, and the path of totality is extremely narrow, so as seen from any one location, or even any one country, a total eclipse is an extremely rare event. When it happens, however, it offers us the only chance of seeing the Sun’s atmo-sphere from ground level without using special equipment. It was through observ-ing eclipses that astronomers in the nineteenth century fi rst began to discover the properties of the Sun’s atmosphere.

The lowest layer of the atmosphere is called the chromosphere , which extends to only a few thousand km above the photosphere. Immediately above the photo-sphere there is a thin layer of gas that is somewhat cooler than the solar surface, but the chromosphere is considerably hotter, with an average temperature of about 10,000 K. It can be seen during a total eclipse as a thin band of light around the silhouette of the Moon, appearing red or pink in color. This color led astronomers to give it its name, chromos being the Greek word for “color.” It glows pink because it is composed mostly of ionized hydrogen that emits light at a distinct series of wavelengths, the brightest being the so-called hydrogen-alpha (H-alpha) line, which glows red. Because it emits at discrete wavelengths, the chromosphere can be viewed from Earth without having to wait for an eclipse, using special instru-ments that block out all other wavelengths. Special fi lters allow even amateur

Fig. 1.2 Total solar eclipse, photographed by Derek Hatch on August 1, 2008, showing the extensive white corona around the black silhouette of the Moon

6 1 Our Sun

astronomers to see this part of the solar atmosphere on any sunny day, as we shall discuss in Chap. 6 . This is very fortunate because, as we shall see shortly, much dramatic activity occurs in the chromosphere.

The main part of the Sun’s atmosphere – and the most striking part during a solar eclipse – is the corona . During an eclipse this is visible as a broad ring of brilliant white light that has no de fi nite outer edge but gradually fades with distance from the Moon’s silhouette. It shines due to scattering of light from the photosphere by protons and electrons. It is a million times fainter than the brilliant photosphere and so is totally invisible from Earth except during a total eclipse. Special fi lters do not reveal it because it shines at all wavelengths and not at certain discrete wavelengths which can be isolated, as the chromosphere does.

The corona can be observed, to some extent, from mountaintop observatories using special instruments known as coronagraphs, but even these reveal only the brighter parts of the corona; the best views of the Sun’s outer atmosphere are obtained from spacecraft.

Between the chromosphere and the corona is a thin layer known as the transition region , in which the temperature increases by a factor of 100. The corona is far hotter even than the chromosphere: its average temperature is around 2 million K. It is so hot that it shines in X-rays and the extreme ultraviolet. This intense radiation is blocked by Earth’s atmosphere – thankfully, because it would quickly kill all living matter on Earth – and so it requires space-based instruments to be observed. Why the corona is so hot is one of the great unanswered questions about the Sun and the subject of much research by professional solar astronomers. They believe that the cause of the heating is connected with the Sun’s powerful, ever-changing magnetic fi eld.

The corona has no de fi nite outer edge, and instead thins out gradually with distance from the Sun. In fact, it continuously emits a stream of protons and elec-trons known as the solar wind . These particles fl y outwards from the Sun at an average speed of 400 km per second and are constantly replaced by new matter. The solar wind pervades the entire Solar System – indeed, it extends outwards well beyond the orbits of the eight major planets, and no one knows precisely where its in fl uence ends and gives way to the much more rare fi ed gas between the stars – though in late 2010 scientists announced that NASA’s Voyager 1 spacecraft, launched in 1977 to explore the outer planets and now headed for interstellar space, had encountered a region where the outward motion of the solar wind had ceased, some 17.3 billion km from the Sun, or approximately 115 times the dis-tance between Earth and the Sun.

Solar Activity

Aristotle and most other western philosophers of ancient times believed the Sun to be a ball of pure white fi re. However, as early as two centuries before Christ, astronomers in China began reporting occasional dark patches on the face of the

7Solar Activity

Sun when it was setting. These were the fi rst reports of sunspots (Fig. 1.3 ), whose presence was con fi rmed by Galileo Galilei and other astronomers in the early seventeenth century using the newly invented telescope.

A sunspot appears dark because it is about 2,000 K cooler than the surrounding photosphere. As the temperature of the photosphere is around 5,800 K, a sunspot is still extremely hot by Earthly standards and only appears dark compared with the brilliant photosphere. If it were possible to isolate a sunspot from its surroundings it would still shine very brightly. Even on the Sun, sunspots are not black but have a brownish tint (Fig. 1.4 ). Sunspots come in all shapes and sizes – from single, isolated spots to complex groups containing anything from a few spots to over a hundred. Sunspots are forming and dying out all the time, and their appearance is never quite the same from one day to the next.

The size of sunspots gives us some appreciation of the vastness of the Sun. A typical small- to medium-sized sunspot is about the same size as Earth, while some of the largest spots have diameters ten times that of Earth.

Sunspots and other solar features slowly move from east to west as the Sun rotates on its axis. The Sun has a mean rotation period, as seen from Earth, of about 27 days, so a sunspot takes just under 14 days to move from the east limb to the west limb. The change in position of the spots due to the Sun’s rotation is very noticeable from one day to the next.

The number of sunspots present is not constant but varies in a cycle from a maximum down to a minimum and back to maximum again over a period of around 11 years. Astronomers call this the sunspot cycle, and the maximum and minimum

Fig. 1.3 The Sun’s photosphere (“surface”), showing a very large sunspot group and several smaller groups, photographed by the author in August 2002

8 1 Our Sun

are known as “sunspot maximum” and “sunspot minimum.” At maximum the Sun is heavily spotted on most days, whereas sunspots may be entirely absent for days at a stretch at minimum. Strictly speaking, the term “sunspot cycle” refers only to sunspots. Other, less visible forms of solar activity also wax and wane over the same period, and when referring to these or solar activity as a whole astronomers use the term solar cycle .

Why solar activity varies in this manner is another great mystery of solar sci-ence; the cause is probably connected with distortions in the Sun’s magnetic fi eld, described below. What is known is that while its period has, on average, been 11 years long since telescopic observations of the Sun began, each cycle has different characteristics. Some cycles are a year or two longer or shorter than others. Some are much more active than others, with more sunspots visible at maximum. Also, the graph of solar activity is different for each cycle. Sometimes the cycle takes longer to reach maximum, and activity can remain the same or even decline for several months before resuming its upward trend. One common feature of all sunspot cycles is that it takes longer to drop to minimum than it does to reach maximum – typically, the rise to maximum takes 3–4 years, whereas it can take 6–7 years to fall back to minimum again.

At the beginning of the twentieth century astronomers discovered that sunspots are strongly magnetic. The magnetic nature of sunspots is an indicator of a fundamental

Fig. 1.4 The planet Mercury in transit across the Sun, photographed by the author on May 7, 2003, with an 80 mm refractor and full-aperture solar fi lter. Note how the silhouette of Mercury, visible near the bottom of the picture, appears jet black, while the small sunspot near the top has a brownish tint, because sunspots appear dark only by contrast

9Solar Activity

property of the Sun: our star has a huge and powerful magnetic fi eld, which astronomers believe is generated by complex motions within the solar interior. Sunspots are believed to be caused by the magnetic fi eld being distorted by the Sun’s rotation. Because the Sun is a gaseous object, not a solid one, its rotational speed is not the same all over. At the equator the solar surface rotates once every 25 days, but near the poles the rotation period increases to 30 days. This differential rotation causes the magnetic fi eld lines to become tangled and gives rise to local-ized “knots” of very intense magnetism. A strong magnetic fi eld stops the passage of energy, and when such a fi eld pierces the Sun’s surface it prevents some of the light and heat from welling up from the solar interior, giving rise to a darker, cooler patch on the surface – i.e., a sunspot. Astronomers think that over a period of years, the differential rotation causes the localized magnetic fi elds to peak in intensity and complexity, thus producing many spots, before “unwinding” into a simpler struc-ture again over the next several years, leaving the Sun with few or no spots. This is thought to be the underlying cause of the solar cycle, though much research remains to be done in order to pin down precisely how it happens.

Sunspots tend to form in pairs, and even large sunspot groups often consist of many small spots and a pair of larger spots. Each sunspot in a pair has an opposite magnetic polarity – one spot negative, the other positive. However, the polarities are reversed in the opposite hemisphere of the Sun. For example, if a sunspot pair in the northern hemisphere has a preceding spot (i.e., the westernmost spot of the pair as the Sun rotates from east to west) of positive polarity and a negative following spot, then in the southern hemisphere it will be the other way around – preceding negative, following positive. These polarities switch at the end of a sunspot cycle, at minimum activity. In the next sunspot cycle, therefore, preceding spots will be negative in the northern hemisphere and positive in the southern hemisphere. The polarities do not return to their original state until the completion of the second cycle, and so the Sun can be said to have a magnetic cycle 22 years long – twice as long as the sunspot cycle.

A good analogy for a sunspot pair is a simple bar magnet, each end of which is of opposite polarity. If you perform the well-known school experiment of placing the magnet under a sheet of paper which has been evenly spread with iron fi lings, the loops of the magnetic fi eld lines become apparent. Sunspots also have magnetic fi eld loops around them. Although they are invisible in ordinary light, they show up in other wavelengths, especially in X-rays and the extreme ultraviolet. Images taken in these wavelengths by spacecraft such as the Japanese Yohkoh satellite, the Solar and Heliospheric Observatory (SOHO), the Transition Region and Coronal Explorer (TRACE) and the newer Solar Dynamics Observatory (SDO), show extremely hot gas fl owing along the magnetic fi eld lines above sunspots. The mag-netic fi elds around sunspots are huge, and, in fact, sunspots are just the tip of the iceberg in a large fi eld of magnetic activity. For this reason, professional astrono-mers refer to sunspot groups as active regions .

In the chromosphere, interaction between magnetic fi elds causes a particularly energetic form of solar activity: solar fl ares . A powerful fl are can release as much energy as billions of hydrogen bomb explosions. All but the strongest fl ares are

10 1 Our Sun

invisible in normal light, but they are very prominent at certain limited wavelengths of the visible spectrum (including H-alpha) and so can be detected from Earth using special instruments. However, most of the energy in fl ares is emitted at X-ray and radio wavelengths, where they can sometimes outshine the entire Sun.

Astronomers believe that fl ares are triggered when complex magnetic fi elds in active regions become wound up in tight spirals and connect with fi eld lines of opposite polarity, causing a “short circuit” effect and therefore an explosive release of energy. Flares most commonly occur above large, complex sunspot groups, where magnetic fi elds of opposing polarities are often found tangled together.

During a total solar eclipse it is usually possible to see one or two pink, fl ame-like structures emanating from the chromosphere. Sometimes, when part of the chromo-sphere is hidden by the Moon, they appear to be attached to the Moon’s silhouette, which led some early observers to believe that they belonged to the Moon! These are, in fact, another form of chromospheric activity, called prominences (Fig. 1.5 ) – large clouds of hydrogen shining at the same H-alpha wavelength as the chromosphere and visible from Earth outside eclipses using special fi lters and other instruments. Some prominences show little change in their appearance over weeks or months, while others can erupt from the Sun and disappear in a matter of hours or even minutes. Both types often show beautiful, intricate structure.

Fig. 1.5 Total solar eclipse, August 1, 2008, showing prominences around the silhouette of the Moon. This is a short exposure of the same eclipse as in Fig. 1.2 ; it allows the prominences to show up without being drowned out by the bright inner corona (Photograph by Derek Hatch)

11Solar Activity

Solar activity also takes the form of large-scale changes in the corona. From time to time apparent gaps known as coronal holes appear in the corona. Coronal holes are areas in which the magnetic fi eld stretches out inde fi nitely into space and does not loop back into the Sun. They are a means by which the Sun emits material out into the Solar System and so they have some effect on Earth (see below). Except during a total solar eclipse, coronal holes and other features of this part of the solar atmosphere are only visible on images taken with X-ray and extreme ultraviolet telescopes above Earth’s atmosphere (Fig. 1.6 ). The familiar photosphere is too cool to emit X-rays and so appears as a dark sphere on such images. On the other hand the corona, at 2 million K, glows brightly at these wavelengths, and any active regions present show up brilliantly against the dark photosphere. Coronal holes appear as dark regions with little or no activity.

An even larger form of solar activity are the huge “bubbles” of coronal gas that spew out of the Sun and become shock waves in the solar wind. These coronal mass ejections (abbreviated to CMEs) (Fig. 1.7 ) are usually connected with erupting prominences and are sometimes associated with fl ares. They broadly follow the solar cycle both in their frequency and their intensity. CMEs occur over a period of several hours, and the clouds of ejected matter can occasionally become larger than the Sun itself before dispersing. At the time of writing, the large-scale structure of

Fig. 1.6 The Sun’s corona, imaged in the extreme ultraviolet by the SOHO satellite in September 1997, showing bright active regions and dark coronal holes. The active regions broadly correspond to sunspot groups in visible light. The photosphere appears dark because it is too cool to emit at this wavelength (Image courtesy of NASA)

12 1 Our Sun

the corona is being more or less continuously monitored by the SOHO satellite and also the STEREO mission, a pair of spacecraft operating simultaneously from dif-ferent positions in Earth’s orbit, one ahead of Earth and one behind, thus allowing the solar wind to be imaged in three dimensions. Time-lapse movies assembled from SOHO and STEREO images of the corona dramatically show the expanding shock waves of CMEs. You can view these movies at various websites, such as http://sohowww.nascom.nasa.gov/ and http://stereo.gsfc.nasa.gov/ .

The Sun’s In fl uence on Earth

Astronomers have long wondered whether the Sun has an effect on Earth’s weather and climate. In recent years, this question has become highly contentious politi-cally, because of concerns about “climate change” – a rapid increase in average global temperatures over the past 50 years (often referred to as “global warming”) and associated changes in weather patterns. Some of these changes can be eco-nomically and socially damaging, especially to many poorer countries. It is still unclear whether these changes have been due to man-made causes, such as increas-ing emissions of carbon dioxide (CO

2 ) and other “greenhouse gases,” or to natural

phenomena, such as solar activity.

Fig. 1.7 A coronal mass ejection (CME) imaged in visible light by the SOHO satellite in November 1997. The Sun’s brilliant disc has been arti fi cially hidden by an occulting disc inside the telescope (Image courtesy of NASA)

13The Sun’s Influence on Earth

After centuries of research and speculation, the jury is still out on the question of how much the Sun affects terrestrial weather and climate. We can certainly say that sunspot maxima and minima do not coincide with warmer or cooler spells on Earth. Nor does a large sunspot cause a period of unusually warm or cool weather. Most scientists now believe that the effects of solar activity are too weak compared with the main forces that drive Earth’s weather to in fl uence weather systems.

There is some evidence that solar activity may affect our climate over longer periods. During the period from about 1645 to 1715 Britain and Europe experi-enced a “Little Ice Age,” in which summers were unusually cool and winters severe, the River Thames, for example, frequently icing over. The study of annual growth rings in tree trunks has shown that trees grew less during this period.

This spell almost exactly coincides with a period of very low sunspot activity in which even at solar maximum very few sunspots were seen. This dearth of solar activity in the seventeenth and early eighteenth centuries was fi rst noticed in the historical records at the beginning of the twentieth century by the British astronomer Edward Maunder, and for this reason it is known as the “Maunder Minimum.” But as to contemporary global warming, the Sun can probably be ruled out as the main cause. Over the past 50 years, the level of solar activity at the maxima of the sunspot cycles has gradually declined from a peak in the 1950s, while global temperatures have continued to rise at an ever-increasing rate. Therefore, the consensus among scientists is that solar activity is at most a secondary cause of climate change.

However, the Sun de fi nitely affects us in a number of other very practical ways. The Sun’s powerful magnetism strongly in fl uences Earth’s own magnetic fi eld, which protects us from many of the potentially dangerous charged particles emitted by the Sun. The solar wind constantly pummels Earth’s magnetic fi eld on Earth’s sunward side. In the vicinity of Earth the solar wind travels at approximately 400 km per second and so compresses the sunward side of the magnetic fi eld. Unable to get any further due to the opposing force exerted by the fi eld, the solar wind particles are de fl ected round the fi eld and eventually fl ow out into interplan-etary space beyond Earth, thus safely bypassing our planet. This “safe haven” in which Earth and its inhabitants are shielded from the continuous stream of solar particles is known as the Earth’s magnetosphere .

During periods of high solar activity, however, the Sun releases intense bursts of particles, often in the form of CMEs, and Earth’s magnetosphere cannot keep them all out. Some particles fl ow down the magnetic fi eld lines towards the magnetic poles and give rise to the aurorae – the well-known “Northern Lights” (aurora borealis) in the northern hemisphere or “Southern Lights” (aurora australis) in the southern hemisphere – which is basically a glow caused by the particles hitting the upper atmosphere. Normally the aurora is only visible towards Earth’s polar regions, but if activity is particularly high a geomagnetic storm can occur, allowing the aurora to be seen at temperate and occasionally even tropical latitudes. A geo-magnetic storm is caused when the Sun releases a particularly large burst of charged particles into the solar wind and this burst reaches Earth’s magnetosphere. Flares, coronal holes and mass ejections can all cause these bursts of particles. A good display of the aurora is a memorable sight, with vivid red and green colors and ghostly patterns of luminous rays and arcs (Fig. 1.8 ).

14 1 Our Sun

The frequency and intensity of geomagnetic storms peaks near the maximum of the solar cycle and bottoms out near minimum. The correlation is far from precise, however. Major storms and auroral displays visible down to temperate latitudes are not unknown near minimum, while at maximum there can be long periods without a large storm. Similarly, the presence of a large sunspot or a large number of sun-spots is not reliable as an indicator of auroral activity. Remember that it is not the sunspots themselves but the invisible activity in the solar atmosphere – i.e., fl ares, coronal holes and mass ejections – that causes the fi reworks.

Some sunspots can cause more numerous and more powerful fl ares than others, while occasionally a really big fl are can come from an innocuous-looking spot. Bursts of particles are not always directed towards Earth and so can miss us entirely. Also, conditions in our own upper atmosphere and magnetosphere can mean that even during a very large geomagnetic storm aurorae are only visible from high latitudes. The only way to be sure of eventually seeing an aurora is to keep checking the sky towards the north each evening. Checking websites such as http://spaceweather.com/ and the Space Weather Prediction Center site, http://www.swpc.noaa.gov/ , may increase your chance of actually seeing an aurora when you look at the northern sky.

Fig. 1.8 A major display of the aurora borealis, photographed from southern England by the author on April 6, 2000. The display was triggered by a major geomagnetic storm

15The Sun’s Influence on Earth

Geomagnetic storms have more effects than colorful displays of the aurora and can affect our daily lives, too. On Earth they can damage power and com-munications systems. For example, the great storm of March 1989, which caused a spectacular auroral display visible from much of the globe, took many people by surprise by introducing a huge amount of electric current into power distribu-tion systems in the United States and Canada. With no warning, several genera-tors burned out, causing six million people to be without electricity for 9 hours. Magnetic storms can also cause major interference in radio communications, with serious consequences for airlines and defense systems. More recently, the massive solar fl ares associated with the large sunspot groups of autumn 2003 (see Fig. 3.10 in Chap. 3 ) caused radio blackouts, and some airplanes fl ying over Earth’s polar regions had to be diverted to avoid dangerous levels of solar radia-tion near the poles.

Also serious are the effects of geomagnetic storms on arti fi cial satellites orbit-ing Earth. Satellites are vital to the modern world in many ways, including naviga-tion, defense, emergency services, weather forecasting, environmental monitoring and telecommunications. Communications satellites are vital in television, the Internet and telephone usage, including mobile phones. If anything happens to a satellite, therefore, the effects could be catastrophic. If a large fl ow of particles enters the magnetosphere during a magnetic storm, a satellite can be overloaded with charge and so burn out, or its on-board computers can reset themselves and so cause malfunctions. An example of such a failure occurred in January 1997 – surprisingly, just a few months after sunspot minimum – when the U.S. communi-cations satellite Telstar 401 underwent a power failure and broke down completely at the same time as particles from a CME hit the magnetosphere. Although we will never be totally sure that the CME was the cause of this particular failure, anoma-lies experienced by another satellite at the same time strongly suggest that the two events were related. Since then, manufacturers of satellites and their on-board instruments have taken care to make them more resistant to geomagnetic storms, though the solar outbursts of late 2003 (mentioned above) caused two Japanese satellites to malfunction.

Solar activity also has a longer-term effect on orbiting space vehicles. During solar maximum, and especially during magnetic storms, the increased numbers of charged particles cause Earth’s atmosphere to expand outwards. This results in many satellites in low Earth orbit losing their orbital speed due to air resistance and thus gradually decreasing in altitude, an effect known as “atmospheric drag.” If a satellite loses too much altitude it eventually slows down too much for it to remain in orbit and succumbs to Earth’s gravity, burning up in the atmosphere as it drops towards the ground. Atmospheric drag caused the Skylab space station to fall to Earth in 1979 and also fi nished off the Solar Maximum Mission satellite 10 years later. Ironically, both spacecraft did important research on the Sun.

Given these potentially serious problems that solar eruptions can cause for human activities on Earth, it is not surprising that both space-based and ground-based solar observatories maintain a constant, 24-hour watch on the Sun and “space weather” – i.e., the solar wind, its associated bursts of particles and their effects on

16 1 Our Sun

Earth’s magnetosphere. NASA’s Stereo spacecraft ( http://stereo.gsfc.nasa.gov/ ) can now monitor CMEs in three dimensions as they stream towards Earth, and it can even image the side of the Sun away from Earth, giving us early warning of solar storms. But this intensive study of the Sun goes far beyond the utilitarian. Large teams of professional astronomers worldwide are attempting to answer the many unsolved problems in solar astronomy, such as why the corona is so hot, what causes the solar cycle and whether and to what extent the Sun in fl uences Earth’s climate. Many are also analyzing solar data in order to develop models allowing us to predict geomagnetic storms and their effects more accurately.

Yet at the same time, the Sun is also accessible with small and simple telescopes, and so amateur astronomers can learn much about the Sun by observing it on a regular basis. It is also possible for the amateur to make useful observations, both in monitoring levels of sunspot activity and in taking meaningful images that can help educate the general public about the Sun.

Now that we have set the scene with this general description of the Sun, the rest of this book describes what the amateur can observe on the Sun, how to record observations and the equipment required to observe the Sun safely.

17L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0_2, © Springer Science+Business Media New York 2012

Equipment for Observing the Sun

Chapter 2

Let us begin this chapter by repeating the warning given at the start of the book. NEVER LOOK DIRECTLY AT THE SUN THROUGH ANY FORM OF OPTICAL EQUIPMENT, EVEN FOR AN INSTANT. A brief glimpse of the Sun through a telescope is enough to cause permanent eye damage, or even blindness. Even looking at the Sun with the naked eye for more than a second or two is not safe. Do not assume that it is safe to look at the Sun through a fi lter, no matter how dark the fi lter appears to be, because the danger of observing the Sun is caused less by its exceedingly bright light – though that would be enough to blind you if you looked at it for long enough – but by its invisible radiation.

The Sun’s Radiation

So we can understand just why looking at the Sun without appropriate protection is dangerous, let us pause to examine the nature of the radiation emitted by our nearest star.

The visible light emitted by the Sun and other stars is just one of many forms of electromagnetic radiation , which is composed of electric and magnetic waves oscillating at 90° to each other. (Electromagnetic radiation also behaves as a stream of particles, but for the purposes of understanding the Sun’s light as it affects the visual observer, it is better to think of it as waves.) Radio waves, infrared radiation, ultraviolet and X-rays are all forms of electromagnetic radiation. Different types of electromagnetic radiation are distinguished by their different wavelengths, which correspond to different energies.

18 2 Equipment for Observing the Sun

γ-rays X-rays UV

Visible Light

Infrared Microwaves Radio

10−11 10−9 10−7 10−5 10−3 10−1 100 101 103

Wavelength / m

Fig. 2.1 The electromagnetic spectrum

The full range of wavelengths when plotted on a diagram is known as the electromagnetic spectrum (Fig. 2.1 ). The shorter the wavelength of electromag-netic radiation, the higher its energy and the greater its danger to humans and other living matter. Thus radiation with the shortest wavelengths and highest energies, X-rays and gamma rays, would quickly kill us if Earth’s atmosphere were not opaque to them and blocked them out. Ultraviolet radiation is partially blocked by Earth’s atmosphere, but some of the longer ultraviolet wavelengths – i.e., those closest to the visible part of the spectrum – reach Earth’s surface. Exposure to ultraviolet light from the Sun can, of course, damage unprotected human skin (and pose a potentially serious health hazard). Most importantly for the solar observer, the atmosphere lets through more than enough ultraviolet to damage the eye, particularly if the Sun is looked at through a telescope.

Visible sunlight is a combination of the seven colors of the familiar rainbow: red, orange, yellow, green, blue, indigo and violet. This visible spectrum comprises a range of wavelengths between 700 nm (1 nanometer = 1 billionth of a meter, or 1 × 10 −9 m) at the red end of the spectrum and 400 nm at the violet end and, as can be seen from Fig. 2.1 , is just a tiny slice of the larger electromagnetic spectrum. Beyond the red end of the spectrum is the infrared. As with ultraviolet, the infrared wavelengths furthest from the visible part of the spectrum are blocked by the atmo-sphere, but those nearest to the visible – called the “near infrared” by astronomers – are transmitted and reach Earth’s surface.

A great deal of the energy emitted by the solar surface is in the form of this invisible infrared radiation, which manifests itself as intense heat and presents the greatest danger of all to the eye. When a telescope is pointed at the Sun, the heat is focused behind the eyepiece and is easily enough to cause paper or clothes to smol-der after a few seconds. When the eye is exposed to the Sun in this manner, the infrared radiation burns the retina (the light-sensitive area at the rear of the human eye), destroying the light receptors. A fi lter may block enough visible light, but it can still pass dangerous amounts of infrared, or ultraviolet, or both without you knowing it. More frighteningly, the retina is not sensitive to pain, and so you may

19Telescopes for Solar Observing

not be aware that your eye is being damaged while you are looking at the Sun. A retinal burn can also spread from a small hole in your vision to a large blind area over the ensuing weeks.

Fortunately, there are nowadays a number of safe ways in which to observe the Sun. Do not attempt solar observing until you have read and carefully considered the advice below. Your eyes are too precious for it to be worth taking risks.

Telescopes for Solar Observing

The fi rst question is, what sort of telescope is best suited to observing the Sun? The short answer is simply to use what you’ve got, as each of the main types of telescope used in amateur astronomy can be employed to some extent for solar observing. We will begin here by outlining the strengths and weaknesses of each type, and so enable you to decide what kinds of solar observing your existing telescope is best adapted to – or allow you to make an informed choice if you are contemplating buying a telescope for solar study. We will then discuss the safety precautions needed to observe the Sun with the various types of telescopes.

Unlike some other branches of astronomy, observing the Sun does not require a large telescope. While many targets in the night sky are dif fi cult to observe because they emit too little light, this is not a problem with the Sun – indeed, the Sun has far too much light! Also, the Sun shows plenty of detail even in a very small tele-scope. Whereas details on planets or deep-sky objects often require a good-sized instrument to be seen, dramatic changes in sunspots and other solar features can be observed and recorded using just a 60 mm (2.4 in.) refractor, provided that it is equipped with appropriate eye protection.

Three basic types of telescopes are used by amateur astronomers: the refractor, which uses lenses to focus the light; the re fl ector, which uses mirrors; and the cata-dioptric, a generic name which covers two different designs combining lenses and mirrors to form an image at the eyepiece.

The Refractor

The refractor (Fig. 2.2 ) is in many ways the best type of telescope for observing the Sun. Perhaps partly because of the recent advent of affordable solar fi lters and the consequent increase in the popularity of solar observing, recent years have seen a great number of new refractors appear on the amateur astronomy market, at all price ranges. All modern refractors use an object glass composed of at least two pieces of glass, in order to eliminate “chromatic aberration” – fringes of bright color around the image caused by the glass splitting up light into its spectrum of colors.

The most common type of modern refractor lens is known as an “achromatic” lens. This type of lens, although greatly reducing the amount of false color, does not eliminate it entirely, and a number of other lens designs have been produced,

20 2 Equipment for Observing the Sun

using exotic types of glass and sometimes involving three or four glass elements (Fig. 2.3 ). Telescopes using these special lenses are known as “apochromatic” refractors, and they give excellent images almost entirely free from false color. They have the advantage of having shorter tubes than achromatic refractors of the same aperture, and this is especially useful in refractors of more than 100 mm aperture, because it makes them more compact and portable. Whereas achromatic refractors generally have focal lengths of between 10 and 15 times their apertures – that is to say, they have focal ratios (abbreviated f/ratios) of between f/10 and f/15 – an apochromat can have an f/ratio of f/8 or even shorter.

Apochromats can be very expensive, particularly in larger sizes, but in recent years some refractors with apochromatic or semi-apochromatic lenses have appeared at more reasonable prices. In apertures smaller than 100 mm an ordinary achromatic

Fig. 2.2 The author’s 80 mm (3.1 in.) refracting telescope on a German equatorial mount with electronic slow motion controls. A home-made balsa wood projection box is attached to the eyepiece end of the telescope for safe viewing of the Sun’s image

21Telescopes for Solar Observing

refractor is adequate for all-around solar observing, provided it has an f/ratio of f/11 or longer; as in small telescopes false color does not hugely degrade the image. Apochromats come into their own in high-resolution imaging of the Sun with webcams or webcam-type cameras designed for high-resolution imaging, because chromatic aberration becomes a serious problem at the very large image scales used in this type of imaging. Do not buy a simple achromatic refractor with a focal ratio shorter than f/10 if you have a serious interest in solar observing. These short-focus refractors are designed for visual deep-sky and comet observing at low powers and are not suitable for the rather higher magni fi cations employed in observing the Sun, because they have far worse optical aberrations than their longer f/ratio cousins.

A small achromatic refractor, with an object glass of between 60 and 100 mm and a focal ratio of between f/10 and f/15 has a number of distinct advantages for the ama-teur solar observer. Firstly, such a telescope is relatively cheap to buy, the prices of 60 mm refractors starting at as little as $100. The second main advantage of a small telescope such as this is portability. Ideally, a telescope used for serious solar observing should be permanently set up in an observatory. However, for it to be worth building an observatory, a site that has access to most of the sky and where objects low on the horizon can be observed is required. Most of us do not have a suitable site, let alone the money to construct a good-quality building. Even if you are lucky enough to have an observatory, the Sun might not always be accessible when you want to observe it. The Sun’s position in the sky varies greatly with the time of day and the time of year. For example, it may still be behind the rooftops in the southeastern sky when you want to have a quick look at the Sun on a winter’s morning before going to work, or it could be equally inaccessible on summer evenings, when the Sun is low in the northwest.

Fig. 2.3 The author’s 60 mm (2.4 in.) f/5.9 Takahashi apochromatic refractor. Its short tube makes it extremely compact and portable

22 2 Equipment for Observing the Sun

For reasons we will discuss in Chap. 3 , there are a lot of advantages in observing the Sun when it is between 15° and 40° above the horizon, and so it is a good idea to have a telescope that is compact and light enough to be carried to a spot where the Sun can be seen. Unlike a larger telescope, a small instru-ment can be carried in one or two pieces and can be used at very short notice – an essential requirement if you live in a cloudy climate, where every minute of clear sky is precious.

Another advantage of a small refractor is that it does not collect much solar heat. This is of great importance if you are using the telescope to project the Sun’s image onto a screen and you are not employing a fi lter over the telescope’s front aperture. When focused inside the telescope, the Sun’s rays generate a great amount of heat, which can cause turbulence in the air inside the tube and so cause the solar image to fl icker and blur. Even more importantly, the concentrated heat inside a medium- or large-sized telescope can be enough to damage your tele-scope’s internal components – particularly if, as with some modern telescopes, the tube contains parts made of plastic. Re fl ectors and catadioptric telescopes both have components close to the focus and so are very susceptible to heat dam-age. Even refractors can collect dangerous amounts of heat, and some of the cheaper modern refractors have plastic drawtubes and other internal parts. Heat damage caused by projecting an un fi ltered image could damage the warranty in some telescopes. But by following the precautions described below in the section on eyepieces for solar observing you should be able to avoid heat problems. Small refractors also allow you to easily attach your own accessories for project-ing the Sun’s image, as will be described below.

Another (albeit expensive) accessory that you might want to attach to your telescope, either now or in the future, is an H-alpha fi lter for observing promi-nences and other features in the solar chromosphere. This is another argument in favor of a refractor, because most H-alpha fi lters are designed to work on refrac-tors, and many adapters are available that enable these fi lters to be mounted on refractors.

If you are considering buying a telescope speci fi cally for solar observing, avoid very cheap models of the kind that are found in camera stores or at the bottom end of telescope manufacturers’ price lists. Such telescopes very often have plastic parts and so can be damaged by solar heat. These telescopes are often bundled with “Go To” systems that fi nd stars and other objects in the night sky at the touch of a few buttons, but are completely unnecessary for locating the Sun. More generally, what-ever kind of telescope you buy, it is best to invest in an instrument with a respected brand name, rather than any of the cheap imitations, called “clones,” that have appeared in recent years. Clones are very tempting, because they appear to offer the same telescope as the brand name for signi fi cantly less money, but they sometimes contain plastic parts, and features such as metalwork, gears and slow motion con-trols are often inferior to the “real thing.” It is better to buy, say, a brand-name 80 mm refractor than a 100 mm clone.

23Telescopes for Solar Observing

The Re fl ector

When we use the term “re fl ector” in amateur astronomy, this most commonly means the Newtonian re fl ector, after its inventor, Isaac Newton. Re fl ectors have often been described as the best all-around telescopes for the amateur astronomer, and with good reason. Of the three main types of telescopes, the Newtonian is the cheapest per unit aperture. Partly because Newtonian telescopes are cheaper, they tend to be sold in larger apertures than refractors. A 114 mm Newtonian is usually the smallest size commercially available, and the 150 mm size has long been popu-lar with the average amateur astronomer. A good 150 mm Newtonian can often outperform a smaller but more expensive refractor in picking up faint star clusters and galaxies and can give very satisfactory results on the Sun, too.

However re fl ectors, especially in the larger sizes, are less convenient than refrac-tors for visual solar observing, for several reasons. The larger size of a re fl ector makes it less portable, and the large aperture is of no advantage in solar observing. Although in theory a larger telescope should resolve more fi ne detail on the Sun than a smaller instrument, the atmospheric turbulence caused by the Sun heating the ground during the day rarely permits this. In an average climate, most of the time you will do well to resolve details as small as 1 arc sec (²) across – the resolution limit of a 100 mm telescope. (1 arc sec = 60 arc min (60 ¢ ) = 1 /

3,600 of a degree, or

approximately 1 / 1,800

of the angular diameter of the Sun or Moon.) In fact, if you use the projection method, the large aperture can be a danger, due to the concentrated heat inside the tube, making the telescope effectively a solar furnace. The second-ary mirror, or “ fl at,” in a re fl ector is especially susceptible to overheating, since it lies quite close to the focus. Even if the heat inside the tube is not great, it can still be enough for the secondary and its mount to generate a small current of warm air, spoiling the solar image. These tube currents are worsened by the fact that a Newtonian has an open tube, which allows the warm air to mix with the cooler air outside, causing further instability.

You can overcome the heat problem by covering the front of the telescope with a cardboard mask that has a small hole cut in it. For a 150 mm re fl ector, the hole should be about 60 mm in diameter, and it should be positioned away from the center of the tube, so that the secondary mirror or its support vanes do not block the light path. Such an “off-axis mask” greatly reduces the amount of heat in the tube but still allows plenty of detail to be seen on the Sun.

Beware of the aperture stop fi tted to the tube caps of some very cheap, small re fl ectors. These are often “on-axis” – i.e., removing a secondary cap from the main cap leaves a small hole in the center. This is fi ne in a refractor, which has no central obstruction, but in a re fl ector the secondary mirror takes up most of the small aper-ture left by the central hole, greatly reducing the contrast of the image and the amount of light coming through and rendering the instrument almost useless. If you are using a 150 mm telescope at full aperture, or a larger re fl ector even with an off-axis

24 2 Equipment for Observing the Sun

mask, you must use a fi lter in front of the telescope, as such large apertures generate too much heat inside the tube for safety.

Although Newtonians are of little advantage in ordinary visual observing, they can be very useful for imaging the Sun at very high resolution using webcams or purpose-built astronomical video cameras based on the webcam principle. Because webcams can capture fl eeting moments when the atmosphere is steady, they some-times allow the large aperture of a re fl ector to be exploited to its full advantage. Webcam imaging of the Sun is discussed in more detail in Chap. 8 .

Catadioptric Telescopes

Since they fi rst appeared on the amateur telescope market in the 1970s, catadioptric telescopes have become very popular with amateur astronomers. Two designs of catadioptric telescope are readily available: the Schmidt-Cassegrain telescope (commonly abbreviated to SCT) and the Maksutov. Both designs work on the same general principle: the light passes through a specially designed corrector lens at the front of the tube before the main mirror re fl ects it onto a much smaller convex mir-ror set just inside the corrector lens. The convex mirror then re fl ects the light back down the tube through a hole in the main mirror, and the image is focused and viewed at the bottom of the tube, as in a refractor. The main physical difference between the two designs is the shape of the corrector lens.

Because they re fl ect the light twice along the length of the tube, SCTs and Maksutov telescopes fold up a long focal length into a short tube, and so they are very compact and lightweight. This makes them very portable and also easy to mount rigidly. The latter quality means that SCTs and Maksutovs are well-suited to astronomical imaging and many of them are marketed for this purpose as well as visual observing. Telescopes of this type designed for serious imaging of the night sky are available in apertures of 200 mm and above, but catadioptrics designed for visual observing and short-exposure photography are also available in smaller aper-tures. Widely advertised models of these small telescopes include the Meade ETX range of Maksutovs (available in apertures from 90 mm up to 125 mm – see Fig. 2.4 ), the exquisite and expensive 90 mm Maksutov made by Questar and vari-ous Schmidt-Cassegrain and Maksutov instruments offered by Celestron.

All catadioptric telescopes can be used for solar observing, but they must be used with a fi lter. These telescopes contain some delicate internal components, and the Sun’s heat could easily damage them if they were used for projection. In any case, the mountings of these telescopes would make it dif fi cult to attach projection apparatus to them. Fortunately, many front-aperture fi lters are available to fi t speci fi c models of SCT and Maksutov telescopes. These instruments are also especially suitable for solar photography and webcam imaging, as they were designed with astronomical imaging in mind, and many adapters and attachments are available for them.

25Telescope Mountings

Telescope Mountings

Just as important as the telescope itself is the mount that supports it and moves it to track the apparent motion of the Sun. A telescope on an inadequate mount is very frustrating to use, as the image constantly dances about when the telescope is adjusted or focused. Even wind or passing traf fi c can cause annoying vibrations in a telescope that is on too fl imsy a mount. Many very cheap, small telescopes sold in department stores and smaller shops are on very poor-quality mountings, and these telescopes – which are often optically inferior anyway – should be avoided. Even some more expensive telescopes of high optical quality come with inadequate mountings, which is why some serious observers re-mount their telescopes on more expensive, heavy-duty mounts.

Mountings for both amateur and professional telescopes are of two basic types – alt-azimuth and equatorial . The simplest is the alt-azimuth, in which the tele-scope moves around horizontally (azimuth) and up and down vertically (altitude). An equatorial mounting is effectively an alt-azimuth mounting that has been tipped over so that the azimuth axis points towards the celestial pole.

Alt-azimuth mountings for small refractors are on tripods, and the telescope is carried on either a small fork or a pan-and-tilt head similar to that on a camera tripod. Alt-azimuth re fl ectors are usually on “Dobsonian” mountings,

Fig. 2.4 A Meade ETX 90 mm (3.5 in.) Maksutov telescope equipped with a glass aperture fi lter. Note that the fi nderscope has been removed. This is an important safety precaution, because the Sun is just as dangerous to look at through the fi nderscope as the main telescope, and careless people, particularly children, could look through it by accident

26 2 Equipment for Observing the Sun

a form of alt-azimuth mount that runs on Te fl on bearings and has a low center of gravity, giving a very stable mount with smooth motions for a re fl ecting telescope’s relatively bulky tube. Equatorial refractors and Newtonians nor-mally employ the “German” equatorial mount design (see Fig. 2.2 ). One disad-vantage of the German equatorial is that it requires a counterweight to balance the telescope, thus adding to the weight of the overall setup. Catadioptric tele-scopes also sometimes have German equatorial mounts, but more usually they are on “fork” mountings, an equatorial design suitable only for telescopes with short tubes but which dispenses with the counterweight, making such tele-scopes even more portable.

An alt-azimuth mount is the cheapest option, and for a refractor or Newtonian owner it is also the most portable. General solar viewing, sunspot counting and even basic drawings and photography can be done with a telescope on a good alt-azimuth mount. As you become more experienced, however, you may wish to upgrade to an equatorial mount. The main disadvantage with an alt-azimuth mount is that in order to keep the Sun centered in the fi eld of view you need to move the telescope in two directions – altitude and azimuth – using two slow motions, which is tedious and compromises accuracy if you are recording sunspot positions or doing detailed sunspot counts. An equatorial mount allows the Sun to be tracked with only one motion.

Best of all is an equatorial mount with a motor drive to keep the Sun in the fi eld of view. Also, an equatorial mount that is properly aligned with the pole enables the orientation of the Sun’s image to remain the same throughout the period in which it is observed – an essential feature for serious sunspot counting and draw-ing the Sun’s disc to determine the positions of sunspots. With an alt-azimuth mount you need to adjust the orientation every few minutes to ensure accurate results. Electronic slow motions are a very useful accessory for a motor drive, as they allow you to scan the Sun at high magni fi cation and make fi ne adjustments to the tracking.

Some modern telescopes have alt-azimuth mounts with motorized tracking, the positional adjustments being made by the instrument’s internal computer. However, although these mounts keep the Sun conveniently in the fi eld of view, they do not correct for the changing orientation of the image, and so you still need to adjust the image to keep it aligned. In any case, computer-controlled alt-azimuth mounts are generally sold with “Go To” systems, which, as we noted above, are not necessary for fi nding the Sun.

Whatever type of telescope mount you choose, make sure it is steady enough. A good test of a mount’s stability is to give the telescope tripod a sharp knock and time how long the vibrations in the image take to die out when a high magni fi cation is being used. If the vibrations take longer than a few seconds to settle down, you may wish to consider upgrading to a better-quality mount. Similarly, check that the slow motions – manual or electronic – turn smoothly in all directions.

27Viewing the Sun’s Image

Viewing the Sun’s Image

Solar Projection

Solar projection is traditionally the most popular method of viewing the Sun’s image. Reduced to its basics, projection consists of holding a simple white screen, such as a sheet of white paper, 30 cm or more behind the eyepiece of the telescope and focusing the telescope so that a sharp image of the Sun is formed on the paper. The image on the paper is formed in a similar manner to that produced on a digital sensor by a camera lens.

In the opinion of many, projection remains much the best method for observing the Sun with a refractor. It is less convenient with a re fl ector, however, because a projection screen attached to the side of the telescope tube is ungainly and can cause problems with the telescope’s balance. Projection is totally unsuitable for Maksutov or Schmidt-Cassegrain telescopes; if you have one of these instruments, you need to use a fi lter over the telescope aperture, as described below.

As an experiment, you can project the Sun’s image quite simply by holding a sheet of A4-sized white paper between 30 and 50 cm behind the eyepiece of your telescope (Fig. 2.5 ). Shade the image from the ambient sunlight by fi tting a sheet

Fig. 2.5 Projecting the Sun’s image onto paper for safe solar viewing. A card has been fi tted over the eyepiece end of the telescope to shade the image from direct sunlight

28 2 Equipment for Observing the Sun

of paper around the telescope tube. A low-power eyepiece will project an image of the Sun several centimeters across – easily large enough to show the principal sunspots visible on the disc. Eyepieces of higher magni fi cation will show more detail in the sunspots but give a fainter image. Moving the paper further behind the eyepiece also gives a larger but fainter image.

Some small refractors, and occasionally Newtonians, come with a ready-made solar projection screen. This usually takes the form of a small sheet of white-painted metal attached to a long stalk, the other end of which is clamped to the eyepiece end of the telescope. This setup will also show the main sunspots. However, while commercial screens such as this and the simple method above are fi ne for experimentation or casual viewing, they are not adequate for serious observing. The main reason is lack of contrast. Even with a sunshade, the daylight makes the image much fainter than it should be and can all but wash out delicate solar details. Also, you cannot make accurate observations with a hand-held or commercial projection screen.

Fortunately, it is possible to shade the Sun’s image very effectively using a box that holds the screen to the telescope while keeping out most of the daylight. Projection boxes are not available commercially, but it is quite easy to make one yourself. The entire box is covered, except for one side, all or part of which is kept open to allow the Sun’s image to be seen. Attaching the box to the telescope leaves both your hands free to make notes, drawings and adjustments. The telescope eye-piece protrudes into the box through a hole in the eyepiece end of the box just big enough to fi t over the eyepiece or drawtube and projects the image onto a white screen at the far end of the box.

In order not to upset the balance of the telescope tube, you need to make the box from a very light material. Balsa wood is excellent, as it is light but strong and can easily be cut and manipulated using ordinary hand tools. You can obtain balsa wood from toy shops, model shops or suppliers to the packaging industry. How you build your box is up to you, as amateurs have achieved excellent results with many dif-ferent designs (Fig. 2.6 ). The author built his own projection box (Fig. 2.7 ) using balsa wood, and the details of its construction are described in Appendix A.

A few amateurs, particularly those with observatories or who do their observing from indoors, mount their projection screens entirely separate from the telescope, on a tripod or other support. This is not a recommended practice for accurate work, as even if the telescope has a motor drive the Sun is always moving relative to the screen. However, observing from indoors out of an open window has much to be said for it, as your surroundings are generally darker, and so the image contrast is increased still further. It is only possible in the warmer months, though, as in winter the mixture of warm and cold air causes severe air currents that ruin the image.

You may ask why, when a number of solar fi lters have been available commer-cially for years and have been tried and tested to be safe, the old-fashioned projec-tion method is still recommended. The reason is that projection has several distinct advantages over fi lters. The fi rst is that it is very safe. Filters can be damaged, and even small scratches or pinholes in their delicate metallic coatings can allow dan-gerous radiation to pass through. It is also possible for fi lters to be blown off the telescope by gusts of wind or fall off if not properly secured – with disastrous

29Viewing the Sun’s Image

Fig. 2.6 A solar projection screen, designed by Roderick Willstrop of Cambridge University, attached to the 200 mm (8-in.) Thorrowgood refractor at the Institute of Astronomy in Cambridge, England. Here the screen is made from a light framework of plywood – heavier than balsa wood – attached to a large and heavy telescope

Fig. 2.7 A balsa wood projection box made by the author and attached to an 80 mm (3.1-in.) refractor

30 2 Equipment for Observing the Sun

consequences for the observer’s eyesight. Projection is also much the cheapest method of viewing the Sun. The projection box described in Appendix A cost less than $10 to build, whereas some fi lters for the same size of telescope cost more than ten times this amount. Thirdly, projection allows the Sun to be observed by more than one person at the same time – a unique advantage over any other astro-nomical object, where people have to take turns at the eyepiece and the focus has to be adjusted each time, or else have to be content with an electronic reproduction of the image on a computer screen. This makes projection ideal for teaching pur-poses and group viewing at astronomy clubs. Finally, for the serious solar observer, a projected image makes it much easier to plot the positions of sunspots and to make accurate counts of sunspot numbers.

Although the projection method described above is generally “very safe,” one word of warning is due. When you use the projection method, the sunlight coming through the eyepiece is totally un fi ltered, and it is only too easy for uninformed people, especially children, to look through the eyepiece by accident. Projection also presents a minor fi re hazard in that you can easily singe hair, clothes or per-sonal items if they are placed too close to the eyepiece. A projection box is a helpful safety device, as it can prevent someone’s eye from getting near to the eyepiece, but the best precaution when using the projection method – as with any other way of observing the Sun – is to never to leave a telescope unattended in the sunshine.

As far as the serious solar observer is concerned, the main problem with projec-tion is the Sun’s heat. Because the telescope is unshielded, heat builds up inside the telescope tube when the instrument is pointed at the Sun. In particular, the eyepiece, because it lies near the telescope’s prime focus, can become very hot, and so great care must be taken when choosing eyepieces for solar projection. The advice of experienced solar observers for many years has been never to use any eyepieces except the very simple types such as the Huygens and the Ramsden, and avoid altogether the more complex types, such as Orthoscopics and Plössls. The reason is that the more complex eyepieces contain several lens elements bonded together with optical cement. The intense solar heat can melt the cement and so damage these eyepieces. Huygens and Ramsden eyepieces each contain just two completely separate lenses, with no cement, and so they can withstand prolonged exposure to the Sun.

Nowadays, however, this advice requires serious quali fi cation. Huygens and Ramsden eyepieces are becoming dif fi cult to get hold of lately, and those that are available are often made of plastic, which will quickly melt when exposed to the Sun. This author personally uses Orthoscopic and Plössl eyepieces for all solar work with the 80 mm refractor, because they give much better images than the simpler types. But do take the precaution of observing the Sun when it is less than 40° above the horizon, when its radiation is much less intense than at higher alti-tudes. For reasons described in the next chapter, it is never a good idea to observe the Sun when it is high in the sky anyway. Also, cover the telescope’s object glass to let the interior of the tube cool after half an hour of continuous solar observation. Provided the above precautions are taken, Orthoscopic or other cemented eyepieces are quite suitable for solar projection with 80 mm or smaller refractors. If you use

31Viewing the Sun’s Image

a re fl ector or a larger refractor, though, you need to reduce the telescope aperture to 80 mm or smaller, using an aperture mask described above.

Do not use “exotic” wide- fi eld eyepieces, such as the Ultra Wide Angle, Nagler or Ethos varieties, for solar projection. These eyepieces were designed for deep-sky observing, and their wide fi elds are not necessary for solar work. They would also cost a lot of money to replace if they were damaged by the Sun’s heat.

If you use a refractor to project the Sun’s image using a relatively high power (as you may need to do when making detailed observations of sunspots), you may notice a small amount of chromatic aberration in the image, which gives sunspots a purplish appearance and can sometimes make it dif fi cult to get a good focus. This can be corrected, to some extent, with a fi lter such as the “Fringe Killer” fi lter sold by Baader Planetarium, which is designed speci fi cally to reduce chromatic aberra-tion. This screws into the front of an eyepiece like a lunar or planetary fi lter, though it is not a solar fi lter as such – it does not block any signi fi cant solar radiation and must be used either with the projection method or in addition to a proper solar fi lter (see below). As when using eyepieces for solar projection, the fi lter should only be used when the Sun is at a relatively low altitude, to avoid overheating. This author has used a Fringe Killer for several years on a 9 mm Orthoscopic eyepiece, which gives 101× with an 80 mm refractor, and it does greatly improve the image at higher magni fi cations. As well as removing a lot of the false color from sunspots and the limb of the Sun, it also enhances the contrast of faculae and especially the solar granulation. At only around $50, it greatly improves a solar image without the need to invest in an expensive apochromatic refractor.

Projection works excellently with refractors, and while it is not so convenient with re fl ectors, it is usually possible to rig up some sort of projection arrangement with this type of telescope as well. You may have to be ingenious and attach some heavy object, such as a small can or bag fi lled with nails, to the bottom end of the tube to balance out the weight of the projection box. But if you have a Maksutov or Schmidt-Cassegrain telescope, or wish to do solar imaging, you have to employ a different method.

Solar Filters

The subject of fi lters is not easy, as many fi lters are not safe for solar viewing. One of the most dangerous types of fi lter is that designed to screw into the eyepiece. These eyepiece fi lters take the form of a piece of dark glass encased in a metal ring and are often marked “SUN.” They sometimes accompany the very cheap, small refractors sold in stores and are another reason for avoiding these telescopes. These fi lters are dangerous partly because they pass unsafe levels of radiation, but also because their proximity to the telescope’s focus can cause them to crack without warning, instantly enabling the Sun’s heat and light to pass unchecked through to the eye. Fortunately, these fi lters are less common than they used to be. Never be tempted to use one, and never use any fi lter at the eyepiece end of the telescope.

32 2 Equipment for Observing the Sun

Only fi lters mounted at the front of the telescope are safe, because they are not subjected to the focused heat inside the telescope and only receive the normal daytime solar heat. Such fi lters are known as aperture fi lters . However, not all aperture fi lters are safe or suitable for solar observing through a telescope. In par-ticular, never use home-made fi lters. They may appear to reduce the Sun’s light and give an image faint enough not to dazzle your eyes, but they let through far too much invisible radiation and so damage your eyes before you realize it. Examples of dangerous home-made fi lters include sunglasses, polarizing and Neutral Density photographic fi lters, CDs, CD-ROMS, DVDs, computer fl oppy disks and food wrappers, including those with a silvery appearance that appear to re fl ect the Sun’s light.

Some older sources recommend using black-and-white photographic fi lm as a solar fi lter. The fi lm is unraveled and exposed to normal daylight before being fully developed, making it turn as black as possible. The silver halide crystals in the fi lm cut out the infrared radiation. However, this is not a good idea for the modern solar observer, as fi lm might still not block enough infrared for use with a telescope. Also, not all black-and-white fi lms contain silver halide nowadays, and in any case, fi lm is becoming ever harder to obtain in the age of digital photography. This advice was intended in the past for those observing a solar eclipse with the naked eye. The optical quality of photographic fi lm is also not nearly good enough for the tele-scopic observer. Even with the naked eye, the Sun has a noticeable “fuzz” around it when viewed through this type of fi lter. Color fi lm of any type should NEVER be used as a solar fi lter, for naked eye or telescopic observing. Although it absorbs visible light well, it is nearly transparent to infrared radiation.

Another traditional fi lter, again intended for naked-eye viewing but worth men-tioning here, is smoked glass – a piece of glass coated in soot by holding it over a candle fl ame. This is very risky because it is easy to coat the glass in too thin a layer of soot, and it is very dif fi cult to get an even layer. There is therefore the risk that the amount of fi ltration might not be enough. The soot layer is also very easy to smudge and damage. In any case, with so many good fi lters available on the market nowadays, there is no need for the modern amateur to resort to such a primitive device.

The only truly safe solar fi lters are fi lters made speci fi cally for solar observing through a telescope. These fall into three basic types. A popular and inexpensive type is a sheet of polyester fi lm, often known by the trade name Mylar, coated with a thin layer of aluminum. Mylar is available from astronomical suppliers either in mounts designed to fi t a particular make and model of telescope or as a sheet from which the amateur can cut a portion to place in a home-made mount. Never confuse silver food wrappers or “space blankets” with Mylar fi lters – they are not produced to optical standards and let through far too much infrared radiation. Some of these materials are not even Mylar but just ordinary plastic coated with silver paint. Before buying Mylar – or, indeed, any solar fi lter – check with an experienced amateur that it is safe and of good quality.

When you take delivery of the fi lter, hold it up to the light and see if there are any scratches or pinholes that let the light through. If it has any such defects, do not

33Viewing the Sun’s Image

use it, as they can let through dangerous radiation. Check with the supplier that the Mylar is coated on both sides; this offers extra protection against scratches or pin-holes. When you receive a mounted Mylar fi lter, you may be surprised that the material is slightly loose and wrinkly in its mount, but this has no effect on its opti-cal performance. In fact, if you choose to mount the Mylar yourself, never pull it taut, as this can damage the surface.

When properly mounted up and attached to the front of the telescope, some Mylar fi lters give a solar image with a strong blue tint. This is because aluminum passes blue light more ef fi ciently than red. It also scatters light slightly, making the sky background look very faintly blue. Both these effects can be reduced somewhat by threading a light yellow fi lter (also available from astronomical suppliers) into the eyepiece, in addition to the Mylar mounted over the aperture. Such “secondary fi ltering” gives the Sun a more realistic color and can also increase the sharpness and contrast of the image.

In the late 1990s, Baader Planetarium introduced a variation on Mylar known as AstroSolar Safety Film (Fig. 2.8 ). In outward appearance it resembles Mylar, but the manufacturer claims that the material has undergone special treatment and that its radiation transmittance has been precisely measured. It certainly gives images far superior to conventional Mylar fi lters. The blue color is hardly noticeable at all – indeed, the Sun looks almost white – and the sky background is almost totally black. The solar image is very crisp, and sunspots show up in beautiful detail. The Baader material is such an improvement over traditional Mylar that it now dominates the market for Mylar-type fi lters. Some companies supply mounted

Fig. 2.8 Aperture fi lter made from Baader AstroSolar Safety Film in a home-made mount, fi tted to an 80 mm refractor

34 2 Equipment for Observing the Sun

fi lters made from AstroSolar, while the material is also available in rolls or A4-sized sheets with instructions on how to make your own mount. AstroSolar is very rea-sonably priced, and buying the material unmounted is even better value for money, as you can share the cost of sheets or rolls with other interested members of your local astronomy club, and you have the option of making up a new fi lter from the remainder of the sheet if your fi rst one gets damaged. Be sure, however, to follow the manufacturer’s instructions for safely mounting the material. Even with a Baader fi lter, a light yellow fi lter will give the Sun a warmer hue and enhance the contrast of the solar image still further. The No. 8 light yellow fi lter is useful in this regard. Secondary fi lters such as this are especially useful in solar photography. Indeed, AstroSolar is also available coated to a lighter density for photographic use, but this is not safe for visual observing. Solar photography with fi lters is discussed in Chap. 7 .

Another solar fi lter that looks like Mylar but isn’t is the R-G Solar Film supplied by Thousand Oaks Optical. This is similarly priced to Baader AstroSolar and is also available either mounted or unmounted. It gives the Sun a yellow-orange tint, much like glass solar fi lters (see below), so a secondary fi lter is not necessary.

If buying a sheet of Mylar-type material, note that its optical quality can some-times vary from place to place on the sheet. It is not uncommon that in one part of a sheet, the solar image is excellent, while in another part of the same sheet it is so fuzzy and distorted as to be unusable, even with the naked eye. It is a good idea to check a sheet of Mylar by looking through it at the Sun or another bright object with the naked eye in order to ensure that you select a sound part of it before mount-ing it up. Also, the quality of Mylar-type fi lters degrades enormously if they are left exposed to air over long periods. A small, sealable plastic bag gives excellent, air-tight protection. The fi lter inside its plastic bag can then be stored in an airtight plastic box of the sort intended for storing foods, in order to protect the somewhat delicate cardboard and plastic structure from knocks and scratches.

A cheaper alternative to Mylar is black polymer, which, although it looks like ordinary plastic, is designed for solar observing and blocks the harmful radia-tion. It gives the Sun an orange tint. Black polymer is somewhat more durable than Mylar, and so is very useful if you travel a lot with your telescope or are going somewhere to see a solar eclipse. Its chief disadvantage is that its optical quality is not as high as that of Mylar or glass solar fi lters, so it does not stand up as well to high magni fi cations in a telescope. Black polymer does, however, make an excellent fi lter for observing the Sun or solar eclipses with the naked eye (see below).

More expensive are the glass solar fi lters supplied by several astronomical instrument makers. These consist of a disc of glass polished on both sides to optical quality and coated on one side with a layer of stainless steel (sometimes known as “Inconel”), chromium or an alloy. Most glass fi lters come in mounts made for a speci fi c telescope (Fig. 2.9 ). They are sometimes available in unmounted format, but the small saving in price is not worthwhile, because glass fi lters are more dif fi cult to mount securely than Mylar fi lters. Again, if you purchase such a fi lter, you should check it carefully for defects.

35Viewing the Sun’s Image

Unlike Mylar, glass fi lters give a solar image colored yellow or orange, due to the different transmission characteristics of the coating, and so there is no need to use a secondary yellow fi lter. The sky background is very dark, and sunspots show up with excellent contrast. Glass photographic fi lters are also available, but again these are not suitable for visual observing. Glass fi lters are more durable than Mylar-type fi lters. Although they can cost up to ten times as much as Mylar fi lters, it may be worth considering purchasing a glass fi lter if you use a catadioptric tele-scope and so have to use a fi lter all the time. Glass fi lters do still need to be stored carefully – use the stiff cardboard box it came in to protect your fi lter from impact and either the plastic sleeve it was originally wrapped in or a sealable plastic bag to keep it from exposure to the air.

Once again, always check out a solar fi lter carefully before using it. Sometimes even aperture fi lters advertised as being suitable for solar observing turn out to be unsafe. A classic case happened in the UK during the early 1990s, when adver-tisements for solar fi lters appeared in a British astronomy magazine. The fi lters were found to be pieces of colored plastic and totally transparent to infrared radiation. One experienced solar observer dramatically demonstrated this at an astronomy meeting by pointing a television remote-control unit through the fi lter and switching on the TV! Enough of the radiation from the unit’s infrared beam was getting through the fi lter to the TV. Only buy fi lters from reputable, estab-lished astronomical suppliers and even then get an experienced observer to check your fi lter before using it.

Fig. 2.9 A glass aperture fi lter mounted to fi t the Meade ETX 90 mm (3.5 in.) Maksutov telescope pictured in Fig. 2.4

36 2 Equipment for Observing the Sun

Other Observing Methods

Before the advent of safe aperture fi lters, there were two alternatives to solar projection. The fi rst was the Herschel wedge , sometimes called a Sun diagonal . This resembles an ordinary star diagonal as supplied with refracting and catadiop-tric telescopes and takes the form of a prism whose housing has a hole at the rear (Fig. 2.10 ). The prism re fl ects a small percentage of the Sun’s light up to the eye-piece while the rest of the light passes through the prism and out of the hole at the back. These devices are not recommended for visual observing, as the amount of radiation passed, including visible light, is still dangerously high, and it is still necessary to use a dark secondary fi lter to dim the Sun’s image enough for it to be comfortable to look at. It is also very easy to burn yourself with the unchecked solar heat emerging from the back of the prism.

In recent years, however, the Herschel wedge has made something of a comeback among more advanced observers imaging the Sun with webcams. The spectacular resolution of these cameras tests aperture fi lters to their limits, expos-ing the tiniest optical fl aws, and so the resolution of the images is sometimes compromised. The high optical quality of the prism in a Herschel wedge allows a good telescope to perform to its theoretical resolution limit, given good enough sky conditions. Also, the very bright image transmitted by a Herschel wedge allows a very short exposure to be used. Using a Herschel wedge with a webcam poses no danger to the eyesight, because the image is focused on a computer

Fig. 2.10 A Herschel wedge or Sun diagonal made by Lunt Solar Systems

37Viewing the Sun’s Image

screen. These devices work best with apochromatic refractors, because they transmit the full range of visible wavelengths (unlike aperture fi lters, which transmit a more restricted range), and so chromatic aberration would show up more with an achro-matic telescope. This characteristic is actually an advantage, because it shows the Sun in its natural color and gives no tint to the image.

Because the Sun’s radiation is coming into the telescope un fi ltered, you need to take the same precautions against heat build-up as you would when projecting the image. For the same reason, Herschel wedges must never be used with catadioptric telescopes or any telescopes with plastic parts inside. Herschel wedges are de fi nitely a tool for the advanced solar imager, something underlined by their high price and the fact that they are usually sold to fi t a 50.8 mm drawtube, a size often found on high-end apochromatic refractors.

Another alternative is the dedicated solar telescope. Several designs for these have been tried, although the best-known is an ordinary Newtonian re fl ector with an non-aluminized mirror. This design works in the same way as the Herschel wedge, the mirror re fl ecting a small amount of light back to the eyepiece and pass-ing most of it out of the back of the tube. Such a telescope still requires a secondary fi lter at the eyepiece and, unless an aluminized mirror were substituted for the unsilvered one, it could not be used for anything but solar observing.

Many astronomy books recommend using a welder’s glass as a solar fi lter. Provided that it is of a shade number 14 or higher (i.e., denser), a welder’s glass is indeed safe for viewing the Sun with the naked eye, as these devices are designed for preventing harmful radiation from reaching the eye. It is not recommended as a telescopic fi lter, however, as its infrared blockage may still not be enough for the far greater power of a telescope. Also, it is not made to the highest optical standards and so would not provide as sharp a view as a conventional solar fi lter. Additionally, a welder’s glass turns the Sun a lurid shade of green.

Observing the Sun with the Naked Eye and Binoculars

For observing the Sun’s larger spots you do not even need a telescope. Very large sunspot groups can sometimes be seen with the naked eye – always provided that it is protected by a safe solar fi lter as described above. A welder’s glass of the appropriate shade number is also suitable as a naked-eye solar fi lter. Black polymer fi lters (Fig. 2.11 , left) are excellent for naked-eye solar observing, but they must be made speci fi cally for solar observing. Also suitable are eclipse glasses (Fig. 2.11 , right), which are intended for observing solar eclipses but can be used for observing naked-eye sunspots at any time.

Never be tempted to examine the Sun through mist or when it is low in the sky – you can never be sure that enough radiation is being blocked out. When you see a large sunspot using your telescope it is interesting to see if it is visible with the protected naked eye, and, as we shall discuss in Chap. 5 , some amateurs carry out systematic counts of naked-eye sunspots.

38 2 Equipment for Observing the Sun

A safe way of observing the Sun without fi lters or optical aid is the “pinhole camera” method. Simply make a small hole no more than 5 mm across in a large piece of cardboard and hold a sheet of white paper 1 or 2 m behind it. This is, effectively, solar projection without a telescope. It is best done indoors, where you can close curtains around the pinhole and increase the contrast of the image. The very poor resolution of this method allows only the largest sunspots to be seen, but it is a traditional, and safe, way of observing the partial phases of a solar eclipse. Remember not to look through the pinhole, or let others do so.

Many astronomy books recommend using a pair of binoculars if you cannot afford a good-quality telescope. This is sound advice and it applies to the Sun as well. The chief disadvantage with binoculars is their relatively low magni fi cation, which allows only the larger spots to be observed well. Binoculars are especially suited to projecting the Sun indoors: simply mount them on a camera tripod and let them throw an image of the Sun on to a white screen mounted 1 or 2 m away. Cover the lenses of one side of the binoculars, or you will end up with two images of the Sun. As with “pinhole camera” solar viewing, closing the curtains around the binoculars will darken the room and so increase the contrast of the projected image. If you are observing outdoors, shade the image from ambient light using cardboard as described above, in the section on solar projection. Be careful not to let the un fi ltered sunlight coming through the binoculars get into your eye, and always keep children or uninformed people under supervision. Also, do not leave binoculars pointed at the Sun for too long, as the concentrated solar heat inside the

Fig. 2.11 A naked-eye solar viewer made from black polymer ( left ) and a pair of Mylar eclipse glasses ( right )

39Viewing the Sun’s Image

instrument may damage the cement holding the prisms in place. You can also use binoculars with a pair of Mylar-type fi lters, provided that you mount the fi lters safely as described above. Some companies supply such fi lters in pairs mounted to fi t larger binoculars. If you have smaller binoculars you will need to buy a sheet of solar fi lter material and make your own mounts.

41L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0_3, © Springer Science+Business Media New York 2012

What Can We See on the Sun?

Chapter 3

Before we discuss what can be seen on the Sun’s surface we need to know what are the best conditions for observing the Sun. Our atmosphere interferes with solar observing in two important respects. First, it produces clouds that obscure the Sun or allow us only a poor view of it. Secondly, there is the constant rippling and blur-ring of the solar image, which is caused by the fact that our atmosphere is in con-stant motion above our heads. Both these atmospheric factors are a nuisance to the amateur solar observer and are the main reason why professional astronomers site their solar observatories on the tops of high mountains or on spacecraft. However, with a little planning and thought about the causes of these problems the amateur can get around them to some extent some of the time.

When to Observe the Sun

Before we discuss the atmospheric dif fi culties there is the question of gaining access to the Sun in the fi rst place. This is not as obvious as it sounds, because the Sun’s position in the sky varies greatly according to the time of day and time of year. As seen from temperate latitudes, the Sun in summer rises on the northeastern horizon and sets in the northwest, passing very high in the south at midday. In winter it rises in the southeast and sets in the southwest, and even at noon it is still quite low in the sky. Overall, the best time of year for observing the Sun is summer, simply because the Sun is then more accessible. Between May and August the Sun is visible until about 9:00 p.m., Summer Time (assuming your local Summer Time is 1 h ahead of local time) and this allows us to observe it during weekdays in the evenings after our work and daily activities are over. In winter solar observing tends to be a weekends-only pursuit for many of us.

42 3 What Can We See on the Sun?

Oddly enough, except in winter, the best time of day in which to observe the Sun is not midday, when the Sun is at its highest. This is because of the phenomenon of convection, in which the strong heat of the midday Sun warms up the ground, which in turn re-radiates this heat into the air, causing bubbles of warm air to rise and “stir up” the atmosphere above. Convection is a major cause of the rippling and blurring of the Sun’s image mentioned above. Astronomers call the amount of rip-pling and blurring the seeing .

In good seeing the rippling is only minor, and the Sun’s image can sometimes appear rock steady for a few seconds at a time. Bad seeing is characterized by a constant, large-scale rippling of the image, with the Sun’s limb appearing to “boil” and fi ne details in sunspots and other solar features becoming invisible. Convection is at its worst around midday, when solar heating is strongest, and for this reason it is wise to avoid observing at this time if you can. This is another reason for choosing summer evenings as a good observing time. By evening the Sun is lower in the sky, its heating effect on the ground is less, and so convection is much reduced.

The early morning is also a good time to observe, because at that time the ground is still cool after the night and the Sun has not yet heated up the ground signi fi cantly. If you can observe at this time, you may get even steadier views, and it is well worth trying. However, for practical reasons the early morning is often an inconvenient time for many people, and good observations are harder to make if you are pressed for time. Many people will still prefer to observe in the evening. The steadiest conditions are often between an hour and 3 h before sunset. Before this time interval the summer Sun is high enough for convection to be a problem. Later on, the Sun is low in the sky, and its light shines through a long path in Earth’s atmosphere, causing the steadiness of the image to decline again.

In winter the solar observer has different problems, as between November and February the Sun never rises high enough to cause much convection. At this time of year observing the Sun at or around midday makes good sense, because the Sun is too low for serious observing in the early morning and late afternoon. But if you live in an urban area, the Sun’s image can be affected by convection from a different source – the rooftops of houses and other buildings warmed by central heating. If you observe the Sun when it is low over a rooftop the seeing can often be ruinous.

When the Sun is very low on the horizon, atmospheric refraction and atmo-spheric dispersion also degrade the image. Refraction is the phenomenon whereby our atmosphere bends the incoming light rays from the Sun, appearing to lift the Sun’s image above the horizon very slightly. The effect varies in inverse proportion to the Sun’s altitude above the horizon, and only becomes noticeable when the Sun is less than 10° high. The effect is made worse because it is not the same over the Sun’s disc. The top of the disc is half a degree higher than the bottom, and so the lower side of the disc is refracted more strongly, causing the Sun’s disc to be dis-torted into an ellipsoidal shape. Obviously this reduces the amount of fi ne detail visible, and it also greatly compromises the positional accuracy of solar drawings and photographs, because the Sun is no longer round.

43Where to Observe the Sun

In dispersion, Earth’s atmosphere acts like a prism when the Sun or a star is at low altitude, splitting the Sun’s light into a tiny spectrum, giving the Sun’s upper limb a blue fringe and the lower limb a red fringe. This effect only becomes noticeable in small telescopes when the Sun is less than 20° high. In larger telescopes it can be corrected, to some extent, with a secondary fi lter that transmits only one color (see Chap. 2 ).

The weather is a major factor in both the accessibility of the Sun and the seeing conditions. Unfortunately for observers in many parts of the world, this is not easy to predict, but some general common-sense rules can be applied. Many amateur astronomers have found that the best seeing usually occurs during high-pressure conditions. In summer these often cause hot, dry, sunny spells, and their winter equivalents cause cold, frosty conditions. A strong area of high pressure can some-times keep other weather systems out, and so such conditions can persist for many days at a stretch. For the solar observer, it is in these conditions that the atmosphere is at its most stable and seeing is consistently good.

Occasionally, seeing conditions can be really superb, the Sun appearing rock steady and details in sunspots and the surface “granulation” (see later in this chap-ter) visible down to the resolution limit of the telescope. Sometimes such exquisite seeing can occur on a sultry and slightly hazy summer evening, when there is little or no wind and the sunlight appears dim and yellow. Conditions such as this are hard to come by in winter, although seeing can be very good on a frosty or slightly foggy day, provided you are not looking over rooftops.

Conversely, bad seeing frequently occurs in low-pressure conditions. After the passage of a cold front, the sky can be beautifully transparent, but the seeing is often disappointing. (Such clear but unsteady conditions can sometimes be good for observing in H-alpha light, when the resolution is less critical but maximum con-trast is required, particularly with some of the cheaper fi lters, which give their best performance in the most transparent conditions – see Chap. 6 .)

Where to Observe the Sun

Your choice of observing site can also have an effect on the seeing, as some atmo-spheric turbulence is caused by local factors. Most of us have very little freedom as to where to observe on a regular basis, as our only site is where we happen to live. However, some locations within a house or garden can provide steadier views than others. We have already noted above that poor seeing results from observing the Sun when it is low over rooftops. Even in summer, when central heating is switched off, rooftops absorb a lot of solar heat during the day, and they constantly re-radiate it back into the atmosphere. The seeing is always far better above trees or grass. For the same reason, seeing can be mediocre in areas containing large stretches of con-crete or tarmac. If you can, try to observe from a lawn or park rather than, say, a driveway or car park – the improvement in seeing is usually quite noticeable. This is one area where a portable telescope scores over a permanently sited one, as it is easy to move a portable instrument to a better site at short notice.

44 3 What Can We See on the Sun?

We noted in Chap. 2 that solar observing does not necessarily have to be done outside. Indeed, much solar observing can be done between May and August from an open fi rst fl oor window that conveniently faces west-northwest, where the Sun is on summer evenings.

Many astronomy books advise strongly against observing the night sky out of a window, and for good reason. At night there is always a large temperature differ-ence between the inside and outside air, even in summer, and opening a window causes warm air to rush out into the colder exterior, resulting in disastrous seeing through the telescope. During the daytime, however, things are different. In the summer months there is often little or no temperature difference, and views out of a window are every bit as good as they are from outside. If you use the projection method, observing out of a window greatly improves the contrast of the solar image, even with a projection box, as the amount of ambient light is greatly reduced, especially if you draw the curtains around the telescope. Observing from indoors has the added advantage of comfort for the observer, as you are shielded from the heat of the summer Sun. Many summer solar observations can be made from an armchair pulled up in front of the projection screen – giving new meaning to the phrase “armchair astronomy”! Finally, if you keep your telescope set up indoors, observing out of a window allows access to the Sun at a moment’s notice – essential if you live in a cloudy climate. During the winter months, though, it is still necessary to observe from outdoors, as any heating inside a room creates enough of a temperature difference to ruin the Sun’s image.

A fi nal source of poor seeing can come from within the telescope itself. Any telescope that is used to project the Sun’s image, or to image the Sun with a Herschel wedge, is a potential solar furnace, because it allows the Sun’s heat to enter the optical system unchecked by any fi lter. The focused heat inside the telescope heats up the air within the tube, causing air currents to circulate. If you observe the Sun’s projected image on a warm day over a period of about half an hour or so, you will fi nd that the image quality gradually deteriorates, regardless of how favorable the weather or your observing site is. The main remedy for this problem is not to observe the Sun for too long continuously, but cover the object glass for a few minutes every quarter of an hour or so to let the telescope cool down.

Heat build-up in the telescope is another reason for avoiding observing when the Sun is at a high altitude, because it is always less of a problem when the Sun is lower in the sky. Another possible solution to the heat problem is to stop down the telescope’s aperture. My 80 mm refractor, like many small refractors nowadays, contains a small secondary cap in its objective lens cover. When the small cap is removed and the main cover left in place, the telescope’s aperture is effectively reduced to about 40 mm. You might sometimes use this arrangement when making your initial reconnaissance of the Sun’s disc at low power, and only remove the main cap when you do sunspot counting and examine details of the spots at high magni fi cation. In this way you can ensure that for some of the time the amount of heat entering the telescope is reduced.

45Aiming the Telescope

Aiming the Telescope

Now that we have decided on the best time and place for solar observing, there is the question of centering the Sun’s image in the telescope. Because you cannot look at the Sun through an un fi ltered telescope, you have to fi nd it by some indirect means. Finding the Sun is a problem whether you use projection or a fi lter, because although an aperture fi lter allows you to view the Sun safely, you still have to aim the telescope at it. You cannot use the fi nderscope, which is every bit as dangerous to look at the Sun through as the main telescope. Neither should you sight the tele-scope on the Sun by looking along the edge of the tube, as this is also dangerous. You could locate the Sun using your telescope’s setting circles and an ephemeris such as the BAA Handbook , and most computerized telescopes include the Sun in their databases. However, this is not an option unless your telescope is permanently sited, as you cannot polar-align your telescope during daylight.

If you have Baader AstroSolar or other Mylar-type material, you could make a small aperture fi lter for the fi nderscope; the only minor inconvenience with this is remembering to install the fi nderscope fi lter each time you observe the Sun. A traditional and safe method of locating the Sun is by using the telescope’s shadow. When a telescope is pointed in the direction of the Sun its shadow is at its smallest. When the telescope’s shadow is as small as you can get it, look at the projected image or, if you use a fi lter, in the eyepiece. You may need to sweep the telescope back and forth slightly to bring the Sun into the fi eld of view. Using a low-power eyepiece makes it easier to fi nd the Sun this way.

Safe as the shadow method is, it also involves an element of trial and error and can be tedious every time you want to observe the Sun. If you live in a damp cli-mate, such as Britain or the northeastern United States, it is often necessary to fi nd the Sun quickly before clouds roll in, and fi nding the Sun by the shadow method can waste several minutes of precious observing time.

Some amateur astronomers, therefore, have made their own “solar fi nderscopes,” which employ the projection method as in the main telescope (Fig. 3.1 ). These take the form of two small pieces of wood or cardboard, one mounted in front of the other along the length of the tube and spaced a few centimeters apart. The front piece has a small hole about 1 or 2 mm in diameter punched in its center. This acts rather like a pinhole camera, throwing a tiny image of the Sun onto the rear piece, which has a cross or mark inscribed on one side. When the miniature solar image falls on the mark the telescope is exactly pointed at the Sun, and the solar image should appear in the eyepiece or projection box. Of course, like a conventional fi nderscope, such a device needs to be accurately aligned with the main telescope in order to work.

Some manufacturers, such as TeleVue and Coronado, sell commercial solar fi nderscopes with designs similar to that described above, but their price tags seem unnecessarily expensive for a simple device that you can make yourself. For my own observing this author uses a variation on the solar fi nderscope (Fig. 3.2 ) on an

46 3 What Can We See on the Sun?

Front disc withhole in centre

Light from Sun

Image of Sun on rear disc,with position marked

Fig. 3.1 Diagram showing the principle of a solar fi nderscope (Courtesy of Dominic Ford)

Fig. 3.2 Solar fi nderscope on the author’s 80 mm refractor. This replaces the “normal” fi nderscope, which has been removed as a safety precaution

80 mm refractor, which is employed mostly for solar observing. The telescope’s conventional fi nderscope, which is of no use for observing the Sun, can be removed. The fi nderscope’s mount is composed of a single metal ring. To each end of the ring you can attach a small disc of Bristol Board (the same thin white card used for projecting the Sun’s image), just big enough to cover the ring. Then punch a small hole in the front disc, using a pair of compasses. The next time you observe the Sun in the telescope using the shadow method, and when its image is accurately centered

47Viewing the Sun’s Surface

in the projection box, carefully mark on the rear disc the spot on which the tiny pinhole solar image fell. On every observing session after that you should be able to locate the Sun instantly by aligning its image with the pencil mark you had made, enabling you to spend more time doing solar observing.

A fi nal point about fi nderscopes. When you are using a telescope to observe the Sun, make sure both lens caps are securely placed on the fi nder – or, best of all, remove the fi nderscope altogether. Lens caps can be removed by curious children, and in any case, it is all too easy to look through a fi nderscope by accident. You have been warned!

Viewing the Sun’s Surface

When observing the Sun, it is a good idea to always start with a fairly low magni fi cation, so that the entire disc of the Sun is visible on the projection screen or in the telescopic fi eld. Not only does this make fi nding the Sun easier, it also enables you to see the entire Sun at a glance and identify any interesting areas. You can then hone in on an interesting sunspot group using a higher power. You can even use two eyepieces with an 80 mm refractor: a 15 mm Plössl (magni fi cation 61×) for viewing the whole disc, and a 9 mm Orthoscopic (101×) for detailed views of solar features.

Usually the most obvious features to be seen on the Sun are sunspots, but before we discuss these let us look at a couple of other general characteristics of the Sun.

Granulation

When viewed at low magni fi cation, the solar surface appears smooth and uniform. However, if you use a medium or high power, and providing the atmosphere is reasonably steady, the Sun shows a faint granular texture, known as the solar granu-lation . The granules are bright polygon-shaped features separated by darker lanes. Each granule is the top of a column of gas rising up from the convective layer inside the Sun discussed in Chap. 1 . The dark lanes are material falling back into the Sun. In the telescope the granules appear tiny, only about 1.5 arc sec across and con-stantly shimmering due to turbulence in our atmosphere. But it gives some indica-tion of the vast size of the Sun when we realize that each granule is about 1,100 km across – larger than the whole of Texas! Granules are very short-lived, lasting only about 20 min on average, and they are constantly fragmenting and re-grouping with other granules. The granulation is not easy to observe in a telescope, partly due to the effects of our atmosphere but also because it is quite subtle and indistinct. These factors make it dif fi cult or impossible to follow the life and death of an individual granule. You may fi nd it easier to see the granulation by slowly moving the tele-scope back and forth using the slow motion controls. Some interference fi lters, such as the “Fringe Killer” supplied by Baader Planetarium, can increase the contrast of

48 3 What Can We See on the Sun?

granulation. The most spectacular views of the granulation are to be had during occasions of exceptionally good seeing, such as the humid summer evenings or frosty winter days mentioned above. In fact, granulation is a good rough measure of the quality of the seeing. If the seeing is very bad, granulation is dif fi cult or impossible to see at all. On most days it is visible to some extent, but turbulence nearly always degrades our view of it.

Limb Darkening

Switching back to a low-power eyepiece, you may notice another characteristic of the Sun’s visible surface: it appears brighter at the centre of the disc than at the edge, or limb. This limb darkening (Fig. 3.3 ) is caused by the fact that the photo-sphere is a thin layer and that its temperature decreases with height, from over 6,000 K at the bottom to just 4,400 K at the top. When we look at the center of the solar disc, we are looking at the photosphere from directly above it, and so we are looking at the hot inner part of the photosphere through only a thin layer of cooler gas. But because the Sun is a sphere, when we look near the limb we are looking at the higher, cooler layers of gas, which appear darker.

Limb darkening has one important effect on solar observing. It allows us to see the faculae , brighter patches on the Sun that often appear around or near sunspot groups. Faculae are usually invisible when they are near the Sun’s center, when they

Fig. 3.3 Whole-disc photograph of the Sun, showing sunspots and limb darkening. The latter has been exaggerated by the high contrast of the image (Photograph by the author)

49Viewing the Sun’s Surface

are masked by the intense brightness of the surrounding photosphere. Faculae are discussed in more detail in a section later in this chapter.

Sunspots

Sunspots are not always present, and within a year or so of solar minimum there may be periods of days or even weeks when there are no sunspots visible. The most recent solar minimum in 2008 was very noticeable in this regard: solar activity remained very low between 2007 and 2010, and over much of this period there were stretches of 90 days or more when there were no sunspots at all. But this minimum, although not unprecedented, was unusually intense and prolonged. Outside the deepest part of solar minimum, there are usually at least a few spots visible on the Sun’s disc.

General Characteristics of Sunspots

The fi rst thing to notice about sunspots is that they come in all shapes and sizes, from the tiniest dots, barely perceptible even when a high magni fi cation is used, to huge, complex blemishes, showing considerable detail even at low power. The smallest spots visible in amateur-sized telescopes have diameters of around 1,000 km, which at the distance of the Sun translates into an angular diameter of 1 arc sec – the resolution limit of a small telescope. A typical small- to medium-sized spot may have a diameter comparable to that of Earth (12,756 km). The very largest sunspot groups can grow to 100,000 km or more in their longest dimension. Large spot groups, however, are not as common as small ones, and the very biggest ones tend to appear only near solar maximum.

You may occasionally see the size of sunspot groups, particularly large ones, being referred to in units called “msh.” This is an abbreviation for millionths of a solar hemisphere and refers to the area of the solar surface covered by the group. Typical, medium-sized groups measure a few hundred msh, while the largest groups can cover over 2,000 msh. The March 2001 group mentioned above was the third largest on record and had a peak surface area of just under 3,000 msh.

Sunspots show obvious changes from one day to the next, in terms of their posi-tion, their shape and their internal structure. In fact, the thought of what might be new on the Sun and what might have changed since yesterday does much to drive the keen solar astronomer to observe the Sun on a daily basis. Outside the pro-longed spotless periods around solar minimum, the Sun’s appearance is never the same from one day to the next.

The main positional change we can see from day to day is that the sunspots move from east to west on the Sun’s disc due to the Sun’s rotation. A typical sunspot takes about 2 weeks to move from the eastern limb to the western limb of the solar disc. As we discussed in Chap. 1 , the Sun’s rotation is differential, i.e.,

50 3 What Can We See on the Sun?

it rotates faster at the equator than it does at the poles. Near the equator the Sun’s rotation period is 25 days, but as we look nearer the poles this increases to 30 days. Long-term observations of sunspots show that spots at high latitudes get “left behind” by those nearer the equator. The effect of differential rotation on sunspot positions is more subtle than the general day-to-day solar rotation, but it is possible to observe it using photography or precise plotting of sunspot positions at daily intervals.

Another interesting thing to notice about sunspot positions is that sunspots are never to be found anywhere near the northern or southern limbs of the Sun. In fact, sunspots are con fi ned to two parallel bands on either side of the equator. Sunspots are not usually seen exactly at the equator, but they can often occur just a few degrees of solar latitude away from it. The main “belts” of sunspot activity lie between 5° and 25° solar latitude – although, as we shall see below, sunspot lati-tudes vary with the solar cycle. Sunspots start to thin out at 30° or so and beyond 40° they are quite rare.

The structural changes in sunspots seen from day to day are many and varied. To give a typical example, a small, single spot seen one day might be a pair of spots the next. A day or two later on it might have grown and contain several more spots. By the end of a week it may have grown to become quite a large and complex group. As a group grows, it tends to fl atten out with respect to solar latitude. For example, a small group may be inclined at a considerable angle to the solar equator to begin with, but as it gets larger this inclination decreases. Very large groups tend to be long and thin, covering many degrees of longitude but often just a few degrees of latitude.

Before we look in detail at sunspots and their characteristics, let us consider a typical sunspot such as that shown in Fig. 3.4 . The photograph shows a simple example, but such sunspots are very common, and the picture illustrates features common to the majority of sunspots. A sunspot contains two sections: a dark cen-tral part, known as the umbra and a somewhat less dark surrounding region, or penumbra . The relative sizes of the two sections vary from sunspot to sunspot, and sometimes a sunspot group contains mostly penumbra and very little umbra; this is often the case with a large group that has passed the peak of its activity and has “decayed.”

At fi rst glance the umbra appears very dark, but it is not really black. As we saw in Chap. 1 , sunspots only appear dark by comparison with the even hotter surround-ing photosphere. Umbrae are, in fact, grayish-brown in color, although in refractors this is sometimes masked by the effect of chromatic aberration, which can give them a purplish tinge.

At fi rst glance, the surrounding penumbra appears smooth and gray through a small telescope, but if you examine the penumbra of a large sunspot using an 80 mm or larger telescope and the seeing is very steady, you may notice some fi ne structure within it. A sunspot penumbra is composed of myriads of extremely fi ne dark streaks radiating outwards from the umbra. These streaks are separated by lighter streaks known as penumbral fi laments (Fig. 3.5 ). These features are only a fraction of an arc second across and so in theory should be invisible in small telescopes.

51Viewing the Sun’s Surface

Fig. 3.4 A small, simple sunspot, showing the dark central umbra and lighter surrounding penumbra (Image by Dave Tyler)

Fig. 3.5 In this remarkably high-resolution image by Dave Tyler of a bipolar sunspot group, taken on April 23, 2011, the penumbral fi laments are clearly visible in two of the group’s spots, especially the large spot on the right. The granulation of the surrounding photosphere also shows up strikingly

52 3 What Can We See on the Sun?

However, high-contrast features with a long, thin shape are often visible beyond the resolution limit of a telescope. The same is true of linear valleys on the surface of the Moon, some of which are only a few hundred meters across. Craters of the same diameter are not visible in the same size telescope.

The type of sunspot described here is known as a “symmetrical” spot, because it has a roughly circular outline when it is near the center of the solar disc. As it approaches the limb, however, the effect of perspective causes it to look elliptical. As long ago as 1769 the astronomer Alexander Wilson of Glasgow University noticed that in a sunspot near the western limb the umbra was displaced towards the center of the disc; in other words, the penumbra on the side of the spot nearest the center of the disc, i.e., the eastern side, was noticeably narrower than that on the side facing the limb. The phenomenon became more noticeable as the spot came nearer the limb and was still visible when the same spot reappeared at the eastern limb a fortnight later, though it was now reversed, with the western side of the penumbra now being the narrowest. The effect became known as the Wilson effect

Fig. 3.6 A small, symmetrical sunspot showing the Wilson Effect: the penumbra on the side of the spot nearest the center of the disc is narrower than that on the side nearest the limb, making the spot look like a depression on the Sun’s surface (Image by Dave Tyler)

53Viewing the Sun’s Surface

(Fig. 3.6 ), and it led Wilson and other astronomers to believe that sunspots were depressions in the solar surface, rather like craters on the Moon.

We now know, of course, that the Sun’s surface is a gas, and that no craters could ever exist there, but Wilson’s conclusions were partly true. Astronomers now believe that the density of the gas inside a sunspot, and especially the umbra, is much lower than that in the surrounding photosphere, allowing us to see about 1,000 km further into it.

However, sunspots do not always show the circular outline described above, with the penumbra concentrically surrounding the umbra. Often the penumbra is more elongated on one side than on the other, and this can cause a false “Wilson effect” when the spot is near the limb. For example, if the penumbra shows more elongation towards the west, this will remain visible as the sunspot approaches the western limb, giving the impression that the spot is a “depression.” Conversely, a spot with a penumbra displaced towards the east can sometimes show the opposite of the Wilson effect: penumbra broader towards the center of the disc, and narrower towards the western limb. For many years astronomers thought that spots with this appearance were small mounds above the solar surface. The Wilson effect is only conspicuous in symmetrical sunspots and relatively simple sunspot groups and is much harder to see in more complex groups.

Evolution of Sunspots

A sunspot begins life as a tiny dot about 3 arc sec or less in diameter known as a pore . Pores are not true sunspots and have a lighter and grayer appearance than genuine sunspot umbrae; as we shall see in Chap. 5 , they are not included in counts of sunspot activity. A pore becomes a sunspot proper when it becomes as dark as a sunspot umbra and increases in size somewhat. Such a spot has no penumbra and borders immediately on granulation on all sides. Small spots such as these are the most common type of sunspots. A spot in its early stages of development is often surrounded by faculae, mentioned above when we discussed limb darkening. Indeed, faculae often herald the formation of sunspots. If you see a prominent clump of faculae at the eastern limb, it would not be surprising to see a small spot in the same region within the next day or two.

Sunspot groups start life either as a single, small spot such as just described, or as a tiny cluster of very small spots or pores. The smallest spots often disappear after a day or less, but if the spot or cluster survives then over the next 2 or 3 days the initial spot may be joined by another spot or cluster of spots close by, the two forming a bipolar sunspot group – that is, a group of two spots or clumps of spots of opposite magnetic polarity. The preceding, or westernmost, spot is known as the leader , and the trailing spot the follower . The separation between the two spots gradually increases. If the group is active, the leader and sometimes the follower will grow and may go on to form penumbrae.

By the end of a week a number of smaller spots may have erupted between the two main spots, the group as a whole perhaps containing ten or more individual small umbrae. By now the group is usually nearing its peak size, and it is common

54 3 What Can We See on the Sun?

for a group at this stage to be dominated by a large leader spot. If the group’s magnetic fi eld is complex, as often happens near solar maximum, the whole complex may grow even larger, and many more small spots may erupt within it. The very largest groups may span more than 100,000 km in their longest dimension and can contain over 100 spots.

Regardless of the size of the group, the larger umbrae, including the leader and follower, often split into two or more sections. The splitting of a large spot is often heralded by a deep notch appearing on one side of the umbra’s longer dimension. Within a day or less, this notch has cut across the spot and split it into two sections. The brighter channel between the two sections is known as a light bridge and is often as bright as the photosphere (Fig. 3.7 ). An occasional variant on the light bridge is a bright “island” of photosphere in the middle of an umbra. If an umbra suddenly splits up into several sections this suggests considerable activity and the possibility of fl ares emanating from the group (see later in this chapter and also Chap. 6 ).

Most sunspot groups reach their maximum size and complexity when they are about a week old. After that time, even a large group starts to decay. The umbrae of the spots get smaller, and more and more of the group’s area is covered with penumbrae. If the group is near the limb at this stage, it will be surrounded by extensive faculae. However, the leader spot tends to remain prominent, with a well-de fi ned umbra and penumbra, even when the rest of the group is dying out.

The longevity of sunspots and groups is as variable as their sizes. The smallest spots and pores last for only a few hours or days, while very large groups can persist for weeks or occasionally months. Because a sunspot takes only about 2 weeks to cross the Sun’s disc, it follows that a long-lived spot can reappear at the eastern limb a fortnight after it has disappeared in the west. This is often the case with a large

Fig. 3.7 A large sunspot group, photographed by the author in July 2004, showing multiple light bridges in the umbra of its leader spot

55Viewing the Sun’s Surface

group of the type mentioned above. When it reappears in the east it has often decayed considerably, the leader and perhaps the follower still being present but only faculae and penumbrae visible between them. Such groups have been known to survive sev-eral rotations of the Sun, becoming smaller each time they come around.

A good example of a large, long-lived group occurred in the spring of 2001. At the end of March the group was a spectacular complex, dominating the fi eld of view at high magni fi cation. It produced some powerful fl ares, and brilliant displays of the aurora were reported from many parts of the world. The group came around again in April, and once more in May, but its size and the number of spots it con-tained were much reduced each time.

Even very large groups are not always long-lived, however. Ten years earlier, for example, again in March, there was a similar huge group that produced fl ares and aurorae, but it did not reappear after it went around the limb of the Sun.

Since early 2011, it has been possible to track active regions on the far side of the Sun online, and so see if a large group is likely to reappear at the eastern limb. This is possible thanks to the STEREO solar mission, in which two spacecraft, originally launched in 2006, monitor the Sun from opposite sides of Earth’s orbit, and so active regions on the Sun’s far side can be observed directly – see http://stereo.gsfc.nasa.gov/ .

Sunspot Classi fi cation and Nomenclature

Sunspots are classi fi ed into types according to the stage they have reached in their evolution. Astronomers classify sunspots using a system of three-letter codes devised by American solar astronomer Patrick S McIntosh. This so-called “McIntosh scheme” is a modi fi cation of the “Zurich” scheme originally devised at the Swiss Federal Observatory in Zurich in the mid-twentieth century, which divided sunspot groups into nine broad classes, lettered A, B, C, D, E, F, G, H, and J. The McIntosh scheme distinguishes between some 60 sunspot types and is now the standard scheme throughout the world for classifying the visual appearance of sunspots. It is useful to know how the scheme works, as it is often used to describe spots and groups in reports and it makes it easier to identify the different types of spots as they appear on the Sun.

The fi rst letter of the code is written in upper case and describes the basic type or class of the spot group:

A Single spot or very small group, with no penumbra B Bipolar group, with no penumbrae C Bipolar group. One spot (usually the leader) has a penumbra D Bipolar group up to 10° long with penumbra around both leader and follower E Group between 10° and 15° long with penumbrae around both leader and

follower and many small spots between leader and follower F Group more than 15° long with penumbrae around both leader and follower H Single spot with penumbra

56 3 What Can We See on the Sun?

The second and third letters are both written in lower case and in italics. The second describes the penumbra of the group’s largest spot:

x No penumbra r Penumbra partly surrounds largest spot. Penumbra is “rudimentary” – i.e., not fully

formed and with a granular appearance s Small, symmetrical penumbra, north-south extent less than 2½°. Umbra within the

spot is small and symmetrical a Small, non-symmetrical penumbra, north-south extent less than 2½°. Umbrae within

the spot are separated h Large, symmetrical penumbra, north-south extent more than 2½° k Large, non-symmetrical penumbra, north-south extent more than 2½°. Sometimes

contains spots of opposite magnetic polarity, leading to an increased chance of fl ares (see section on Flares below)

The third letter describes how tightly grouped the spots are in the central part of the group, between the leader and the follower:

x No spots between leader and follower. Applies only to types A and H, where there is only one spot in the group

o Few or no spots between leader and follower. What spots exist are very small i Numerous spots between leader and follower, none of them with fully fl edged

penumbra c Many good-sized spots between leader and follower, at least one of which has a well-

developed penumbra. Again, as with the k category under the leader penumbra classi fi cation above, spots of opposite magnetic polarities are possible in this classi fi cation, leading to an increased likelihood of fl ares

The fi rst letter describes the stage the group has reached in its development. As we have already seen, when a group starts life by turning from a pore into a sunspot proper it is a single spot without a penumbra – i.e., class A. Over the next few days it may become a small bipolar group without any penumbrae (class B). If the leader and then the follower develop penumbrae it becomes class C and then D respec-tively. If the group grows and becomes more complex after this point it may become an E or even an F group. As the group decays and leaves behind only the leader spot, it becomes class H.

Some examples of sunspot group types and their McIntosh classi fi cations are illustrated in Fig. 3.8 . The most common type of group is also the simplest – type A xx , a single spot with no penumbra and, by de fi nition, no spots between the leader and follower. A more complex example is type D ki – a bipolar group with penum-brae around leader and follower, the main spot having a large, asymmetric penumbra and numerous small spots (without penumbrae) between the two main spots. One of the most spectacular types is F kc – a group more than 15° long with a large, asymmetric penumbra around its main spot and the area between the two main spots showing many substantial spots, at least one of which has a well-developed

57Viewing the Sun’s Surface

penumbra. The very large group seen in March 2001 and mentioned above under “Evolution of Sunspots” was of this type (Fig. 3.9 ).

Another giant F kc group appeared in October 2003 (Fig. 3.10 ); this produced many powerful fl ares, including the strongest X-ray fl are ever recorded. Occasionally, in groups of class E kc or F kc , the entire group is enveloped in a single penumbra. Statistical analysis of McIntosh sunspot classes, however, has shown that the major-ity of sunspot groups are of the small and simple types.

Fig. 3.8 Some examples of sunspot group types, ranging from simple to complex, showing their McIntosh classi fi cations. All images by Dave Tyler. ( a ) Very small group of class A xx , ( b ) Small bipolar group, class B ro , ( c ) Spot group of class D ao , ( d ) Spot group of class E kc , ( e ) Symmetrical spot, class H hx

Fig. 3.9 The giant sunspot group of March 2001, McIntosh classi fi cation F kc . This was one of the largest sunspot groups ever recorded and was associated with widespread auroral displays (Photograph by the author)

Fig. 3.10 Another of the giant sunspot groups of Solar Cycle 23, this time the largest of the groups of October 2003, McIntosh classi fi cation F kc (Photograph by the author)

59Viewing the Sun’s Surface

Note that although the McIntosh scheme covers some 60 separate sunspot group classes, not all combinations of letters are possible. For example, if we consider the descriptions above, we can see that, say, types F xx or A ki would be logically, as well as physically, impossible!

When an active region appears on the Sun, it is given a reference number. For example, the large group seen in March 2011 was labeled AR11164, AR standing for “Active Region.” Each active region is issued its own number, which follows in sequence from the previous one. For example, the fi rst active region to be seen after AR11164 was numbered AR11165. The system was introduced by the U.S. National Oceanic and Atmospheric Administration (NOAA) and is now recognized across the world as a standard nomenclature for solar active regions. After you have made a solar observation, it is interesting to look up the AR numbers of the groups you have noted, sketched or photographed, and it is useful to refer to groups by their AR numbers when submitting your reports to solar observing organizations.

Faculae

If you look at a sunspot or sunspot group close to the Sun’s limb, you may notice that it is surrounded by a slightly brighter area or a number of brighter patches. These are faculae (Latin for “little torches”), which we have already noted above under limb darkening and sunspot evolution. They are areas of gas about 300 K hotter than the photosphere, forming a little higher up than sunspots. Nearly all sunspots have at least some faculae associated with them, but the faculae are usu-ally impossible to see when they are near the center of the solar disc and only become distinct when near the limb, where they show up by contrast with the limb darkening (Fig. 3.11 ).

You may notice that some areas of faculae contain no sunspots and appear to exist in complete isolation from any spots or groups. Actually, however, faculae are nearly always associated with sunspots. As mentioned above under evolution of sunspots, an area of faculae often forms before a sunspot or group of sunspots appears in the same position. Similarly, faculae may persist for many days or weeks after a group has decayed and disappeared. With the notable exception of one type of faculae – the “polar faculae” (see below), all faculae form in the same zones of solar latitude as sunspots – i.e., two bands parallel to the equator, stretching to 40° or so. In fact, astronomers believe that faculae are regions produced by magnetic fi elds too weak to cause sunspots.

Like sunspots, faculae come in a large variety of shapes and sizes. They can range from small, bright “blobs” to large, extended patches, while some can appear as long, thin streaks. Some extensive areas of faculae, as are often found around active sunspot groups, have exceedingly complex arrangements. Their low contrast can make them dif fi cult to observe, but one property can help them stand out better.

Faculae emit very strongly towards the blue end of the visible spectrum. As noted in Chap. 2 , Mylar-type fi lters transmit strongly in the blue. This gives the Sun a blue tint as seen through the eyepiece, but many observers have noted that

60 3 What Can We See on the Sun?

faculae appear more distinct through these fi lters than with glass fi lters. If you project the Sun’s image, another way to increase the contrast of faculae is to proj-ect onto a type of white card known as “Bristol Board” (see Appendix A of this book on building a projection box), as this has higher blue re fl ectivity than ordi-nary white paper or card.

Polar faculae form at solar latitudes higher than about 50° and most frequently between 75° and 85°. They are not nearly as easy to see as normal, spot-related faculae. Unlike the latter they do not occur as extended patches but as tiny, isolated bright spots a few arc seconds across, often about the same size as individual pieces

Fig. 3.11 A sunspot group near the limb, showing extensive faculae, imaged by Dave Tyler on March 7, 2011. An image of the same group in H-alpha light, shown for comparison, appears at top

61Viewing the Sun’s Surface

of granulation, therefore requiring good seeing to be visible. They are also very short-lived, never lasting for more than a few days and often with lifetimes mea-sured in minutes. Adding to the dif fi culty of observing them is the fact that, as described in Chap. 4 , the position angle of the Sun’s poles varies, so that during the early part of the year the south pole is tipped towards us, whereas in the later months of the year the north pole is favorably positioned. It follows that faculae are most easily observed around the south solar pole in late winter and spring and around the north pole in late summer and autumn.

Polar facula activity varies in reverse to the activity of sunspots and their associ-ated faculae, in that polar faculae occur in greatest numbers during solar minimum and the years immediately preceding it. During the rise back to maximum they disappear and are relatively rare at maximum.

Polar faculae have not been so thoroughly studied as other features of the Sun, and systematic investigation of them is useful work for the amateur. Although they are much harder to see and keep track of than sunspots and sunspot-related faculae, they can often be seen with just a 60 mm refractor and the projection method with a well-darkened projection box. Using a relatively high power (100×) and waiting for at least reasonably good seeing will increase your chances of fi nding them.

Flares

A solar fl are is by far the rarest and most exotic solar phenomenon you can observe with your telescope in ordinary visible light. As noted in Chap. 1 , fl ares are extremely power-ful in radio and X-ray wavelengths, where they can sometimes outshine the entire Sun for short periods, but they emit relatively little energy in the visible spectrum.

Flares can be detected in certain narrowly de fi ned wavelengths of visible light, using special fi lters that isolate these wavelengths from the rest of the spectrum (see Chap. 6 ), but in everyday white or “integrated” light all but the most intense fl ares are invisible. When a fl are is visible in white light, it is known as a white-light fl are , and if you ever see one it is a major event that needs to be reported quickly. More details on reporting are given in Chap. 5 .

In fact, the fi rst solar fl are to be discovered was a white-light fl are. It was observed on the morning of September 1, 1859, long before instruments for observ-ing the Sun in radio, X-rays or narrowband visible wavelengths had been invented. Two British amateur astronomers, Richard Carrington and Richard Hodgson, were observing the Sun independently of each other when they noticed two brilliant patches of light suddenly appear in the midst of a large sunspot group. The phenom-enon faded away and disappeared from view in a matter of 5 min. The following day, early magnetic instruments recorded a serious disturbance in Earth’s magnetic fi eld, and on the night of September 2 there was also a brilliant display of the aurora. Some astronomers immediately made the connection between solar activity and magnetic storms on Earth, although it was not until near the end of the nine-teenth century that the solar origin of aurorae was widely accepted.

62 3 What Can We See on the Sun?

White-light fl ares are extremely rare. Only a handful have been recorded since the original Carrington-Hodgson observation of 1859. Seeing one demands a great deal of luck, but some characteristics of white-light fl ares can improve your chances. First, sunspot groups of McIntosh classes E and F are overwhelmingly the most likely to produce fl ares; by comparison, fl ares are very infrequent in classes A to D and H. They are especially likely to occur in groups with a compact structure and containing many spots. Sub-classes ki (largest spot has large, asym-metric penumbra and numerous spots between leader and follower) and kc (same, but spots between leader and follower have penumbrae) are the most fl are-productive. So if you see a group of class E kc or F kc , be especially alert for fl ares – in hydrogen-alpha and calcium-K as well as white light. If you see a very bril-liant fl are in H-alpha or calcium-K, it is worth looking to see if there is any sign of it in white light.

Another way of increasing your chances of seeing a white-light fl are is to moni-tor an especially complex group several times a day. Your chances are much reduced if you observe just once a day, as this provides just a snapshot of the Sun’s behavior. Also, like faculae, white-light fl ares emit much of their visible light towards the blue end of the spectrum, so using a blue-transmitting fi lter may help as well.

White-light fl ares show up as small, bright spots, patches or streaks within or very close to a sunspot group. However, be careful not to confuse a light bridge with a fl are. White-light fl ares usually have lifetimes of less than 10 min before fading away, whereas light bridges retain the same brightness for hours or even days. The bright areas of the photosphere that occasionally appear in the umbrae and penumbrae of sunspots can also be confused with fl ares.

63L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0_4, © Springer Science+Business Media New York 2012

Solar Position Measurements

Chapter 4

An important – and very satisfying – task for the amateur solar observer is recording the Sun’s appearance each day and measuring the positions of sunspots and other features. The cheapest and most convenient way of doing this is to make a drawing of the Sun. This may seem paradoxical in an age when digital imaging can capture the Sun’s image in an instant and record details to a level of accuracy beyond even the best draughtsman.

Photography and CCD imaging are in theory more accurate than drawing, and some amateur astronomers use their images to derive sunspot positions; methods of orienting the camera and obtaining high-quality images for measurement purposes are described in Chap. 7 . But using images to obtain sunspot positions is less practi-cal and convenient than drawing, for several reasons. It is not as easy for the aver-age amateur astronomer to de fi ne the cardinal points (north, south, east and west) on a photographic image as accurately as on a good drawing. Secondly, an image taken at a scale small enough to show the whole Sun on one frame – essential for measuring the positions of sunspots – has a relatively limited resolution, and small spots that can be picked up visually can be missed on a photograph. Many digital imaging sensors nowadays have a resolution far higher than even the best 35 mm photographic fi lms, but for positional measurement an image of the whole Sun has to be either printed out or displayed on screen, with an inevitable loss of resolution, unless the screen or print is very large. Whole-disc solar photographs also some-times suffer from distortion towards the edge of the frame, which can compromise the accuracy of sunspot position measurements. Thirdly, a drawing is ready for analysis as soon as it has been completed, whereas images have to be downloaded, processed and printed out. This can be time-consuming, especially if you are sending

64 4 Solar Position Measurements

results in to an amateur observing organization that requires observations on a monthly basis. Finally, drawings continue a historical series of observations going back well over a century, so it is easy to compare present-day observations with those made in the past. At least one professional observatory, Mount Wilson in California, makes a daily drawing of the white-light Sun and publishes it online – see http://www.mtwilson.edu/sci.php .

Drawing Using the Projection Method

You can make accurate solar drawings quite easily using a method employed by amateur solar observers for many years. In your projection box, or on your projec-tion screen (see Chap. 2 and Appendix A of this book), insert a card containing a fi nely ruled grid. Underneath the paper on which you make your drawing is clipped an identical grid, whose lines are thick enough to show through the drawing paper. You can then simply copy the positions of any sunspots present from the grid in the projection box onto their equivalent positions on the grid under the drawing paper, and you have an accurate plot of what is visible on the Sun that day.

This potential for accurate drawings is a major reason why projection is the preferred method over observing by direct vision through a fi lter. Using the latter method, the result is bound to be less accurate, as the drawing has to be done free-hand. You could compromise by using a crosshair eyepiece, such as that used on a fi nderscope, and so divide the Sun’s disc into quadrants. It may also be possible to use a reticule intended for a microscope or a precision measuring instrument, obtained from an optical surplus supplier. But in any case, the projection method is still the most convenient, as it allows you to comfortably look at the Sun’s image with both eyes and involves no eye strain as you switch your eyes back and forth from the image to the drawing.

Making a Projection Grid

Figure 4.1 shows an example of a typical solar drawing grid, consisting of a circle to fi t the Sun’s projected image and divided into small squares. These squares are fur-ther divided by diagonals for more precise plotting. Many observers use a circle 152 mm in diameter for their solar drawings, as this is a standard solar disc size used by many organizations when measuring solar drawings. If you have a very small telescope – say 60 mm or less in aperture – and therefore use a smaller pro-jection box, you may wish to use a 100 mm disc. A 6-in. disc is best divided using Imperial measurements into ½-in. squares. If you prefer to use a 4-in. disc, 4 in. is actually 102 mm, but for accurate plotting you need somewhat smaller squares than the 6-in., and so it is best to use metric units: draw the disc to a diameter of exactly 100 mm and divide it into 10 mm squares.

65Drawing Using the Projection Method

Constructing your own grid is quite easy, although it takes care and patience, since it has to be accurate. In fact, you need to make two grids – one for the tele-scope and one for the drawing. You could draw the grids electronically using draw-ing software and then print them out on a laser printer or a high-quality inkjet, but if you do it this way you need to ensure that the disc diameter and square sizes you have chosen are accurately reproduced – for example, that a 6-in. disc really is exactly 6 in. across. Here we will describe how to construct the grids manually.

To start, you need a ruler and a 360° protractor for the size of disc you wish to use (good stationery shops should supply both 4-in. and 6-in. protractors). The grid to go on the projection screen should be made from a sheet of Bristol Board (the same material as that used for making the screen in a projection box – see Chap. 2 ) or some other form of thin, smooth white card. The grid to be used under the draw-ing can be made from this material or cartridge paper, if you prefer.

Draw everything in pencil fi rst of all, so that you can correct any mistakes. Use a hard pencil that is sharp and stays sharp while you are drawing – a grade 2H or harder pencil is best. Dealing with the projection screen grid fi rst, start by drawing a circle of the appropriate diameter on the white card. Then, using the protractor, mark four points 90° apart around the circumference of the circle. These will later

N

S

EW

Fig. 4.1 A solar drawing grid, consisting of a 152 mm (6 in.) circle divided into ½ in. squares, as used in the projection box and underneath the drawing. The 5° intervals marked out around the circumference are for determining the Sun’s true orientation for sunspot counting purposes (see Chap. 5 ) (Courtesy of Dominic Ford)

66 4 Solar Position Measurements

be the north, south, east and west points when you have worked out which way around the telescope shows the solar image. Join these points to form a cross over the circle. Its center should be at the exact center of the circle.

Then, mark a series of points on the horizontal axis of the cross, half an inch apart for the 6-in. circle, 10 mm apart for the 4-in. Do the same with the vertical axis. Now, move the ruler 30 mm or so above the horizontal axis and again mark a series of points half an inch apart, so that their positions precisely correspond with those on the axis. Do this for each axis.

For even greater accuracy, you could repeat this procedure 30 mm on the other side of each axis. You could use graph paper to help you line everything up accu-rately. Indeed, some observers have used graph paper as an actual projection grid, though the dense network of lines on graph paper can distract from the sunspots and other solar features.

Finally, for each axis, join up the corresponding points with lines stretching across the circle, taking time and care to ensure the lines are precisely parallel. When you have fi nished, you should have a grid composed of squares of the appro-priate size. It is then quite easy to draw in the diagonals, using the corners of the squares as guides. Draw the pencil lines softly so that they are faint, since heavy lines, like graph paper, are a distraction and could cause you to miss some of the smaller sunspots. (If constructing the grids by computer, select a very light gray line color.) Some observers like to letter their squares – for example, alphabetically on the horizontal axis and numerically on the vertical. This helps you to remember on which square you need to plot a sunspot. Otherwise, you can simply establish the spot’s location by counting the number of horizontal and vertical squares from the center of the disc.

You then need to repeat the above process to draw the projection grid that is to go underneath the drawing paper. However, whereas the grid to go on the projection screen should be drawn with faint pencil lines, the drawing grid needs to have black ink lines, so that it shows up clearly through the drawing paper. To ensure the best possible accuracy, the lines should not be too thick. A fi ne-tipped technical drawing pen should produce lines thick enough to be visible through ordinary paper.

In order to orient the image correctly you need to be able to rotate the projection grid. You can do this by rotating the entire projection box, but it may be easier to rotate just the grid. To make the grid rotatable, cut it out to a circular shape and attach it to the screen end of the box with a drawing pin or a small screw through the central axis of the grid (Fig. 4.2 ).

When you have placed the penciled grid in the projection box or screen, adjust the box so that a focused image of the Sun exactly fi ts the 6-in. circle. If you are observing near sunrise or sunset the Sun’s image may appear slightly elliptical, owing to atmospheric refraction (see Chap. 3 ), and so these are not good times to make accurate drawings. The projection screen with its grid needs to be precisely square-on with respect to the telescope’s “optical axis” – i.e., the path a beam of light takes running through the dead center of the tube. If you look carefully at the projected image, preferably with the grid temporarily removed and a plain white

67Drawing Using the Projection Method

screen inserted in its place, you should see at least parts of the edge of the fi eld of view near the corners of the screen. If the edge of the fi eld is concentric with the center of the projection grid and is central within the box (i.e., is at an equal dis-tance from all four corners of the screen), then the box is on-axis. If the screen is notably tilted away from the axis, the Sun’s image is distorted into an ellipse and so the accuracy of any sunspot positions obtained is compromised. You need to look out for this especially vigilantly if your projection box is mounted on its own stand separately from the telescope, for then the telescope is moving all the time with respect to the box.

Orienting the Image

Before you can begin drawing, you need to establish the four points of the compass on the image. Establishing the east-west direction is a simple matter. Using the telescope’s slow-motion controls, position a sunspot on the horizontal line running through the center of the grid. Switch off the telescope’s motor drive (if it has one) and watch the direction in which the spot drifts across the screen. Then rotate the projection grid so that the horizontal line runs along this direction of drift. Repeat this process until the spot drifts precisely along this line. The grid is then oriented east-west. The Sun, like all other astronomical bodies, moves from east to west, so the mark on the circumference of the circle towards which the spot is moving is the west point. You can mark this point “west” for future reference.

Fig. 4.2 Close-up of the author’s solar projection box, showing the rotatable projection grid

68 4 Solar Position Measurements

You then need to establish which end of the grid’s vertical axis is north. If you have an equatorial mount you can do this by turning on the declination’s slow motion so that the telescope moves northwards, towards where Polaris would be if it were visible. If you look at the projection screen, the direction in which the view appears to be moving is north, and you can mark the north point accordingly.

If your telescope is an alt-azimuth, this exercise is best performed near midday, when north is vertically above the Sun in the sky. Whatever type of telescope you use, you will probably notice that north is at the top of the image, in contradiction to the rule that an astronomical telescope gives an upside-down image. This hap-pens because projection involves an extra re fl ection, which “ fl ips” the image north-south. The effect is visible in a refractor or a re fl ector. East and west, however, are not affected, and remain the same as in an astronomical telescope – i.e., west to the left, east to the right.

However, if you are using a refractor and are projecting the image through a star diagonal, east and west are fl ipped, putting east to the left and west to the right, exactly the same as the naked eye view of the Sun. If you are using a Schmidt-Cassegrain or Maksutov telescope, you will be viewing a fi ltered solar image directly through a star diagonal. In this case the image orientation is the same as that of a projected image without a star diagonal – i.e., west left, east right, north top. The diagram in Fig. 4.3 shows the orientation of the Sun’s image as seen in different telescope con fi gurations.

Making the Drawing

You are now ready to begin the drawing. The type of paper you use is a matter of personal preference, but the thinner the paper the better, since this allows the lines to show through more clearly. Tracing paper is ideal, though this is quite expensive if you do many drawings. Some observers use a large, blank notebook for their drawings, with the grid paperclipped behind one of the sheets, while others do their drawings on loose sheets of paper on a clipboard. Whichever method you choose, make sure that both the grid and the drawing paper are steadily mounted and cannot move relative to each other. If you use a clipboard, use a clip at one side of the paper in addition to the main clip at the top of the board. Plot the sun-spots using an HB- or H-grade pencil. Too soft a pencil will quickly wear down and become too blunt to give an accurate rendition, whereas too hard a grade makes it dif fi cult to rub out mistakes (as you will probably want to do quite frequently if you are just beginning!).

Start with the western part of the northern half of the Sun’s image and work eastwards, before switching to the southern half and working westwards. Don’t get bogged down with trying to reproduce the fi nest details in your drawing. The most important aim in a full-disc drawing is to get the positions of the spots correct. Draw in enough detail to show the general shapes of the spots and groups, with the umbrae and penumbrae of the principal spots. If you are drawing a large group, start

69Drawing Using the Projection Method

by establishing the positions of its principal spots on the grid squares and diagonals, and then fi ll in the fi ner details between them, using the main spots as a guide.

If your telescope is on an alt-azimuth mounting, remember that the Sun’s image is slowly rotating, because the mount is not aligned with Earth’s axis. Therefore you need to adjust the orientation every few minutes to keep the east-west line of the grid coincident with the east-west direction on the Sun’s image. This is not a problem with an equatorial mount. Provided that the mount is aligned with the pole reasonably accurately, the orientation should remain the same throughout the obser-vation. Best of all is an equatorial mount with a motor drive, as then you do not have to keep adjusting the RA slow motion control to keep the Sun on the circle.

However, even with a well-aligned equatorial and a good drive, the Sun can still stray slightly off the grid, due to small errors in polar alignment and the

Na

c

e

b

d

f

E

S

W

N

W

S

E

N

E

S

W

S

E

N

W

S

E

N

W

S

W

N

E

Fig. 4.3 Orientation of the Sun’s image at midday as seen in various telescope arrangements. ( a ) Straight projection through a refractor or Newtonian (northern hemisphere). ( b ) Projection using a star diagonal (northern hemisphere). ( c ) Filtered solar image as seen directly in a refractor, Schmidt-Cassegrain or Maksutov telescope, using a star diagonal or 90° mirror (northern hemi-sphere). ( d ) Straight projection through a refractor or Newtonian (southern hemisphere). ( e ) Projection using a star diagonal (southern hemisphere). ( f ) Filtered solar image as seen directly in a refractor, Schmidt-Cassegrain or Maksutov telescope, using a star diagonal. In all arrange-ments with a star diagonal, these diagrams assume that the observer is standing or sitting directly behind the eyepiece and facing the front of the telescope tube (Courtesy of Dominic Ford)

70 4 Solar Position Measurements

drive system. You still, therefore, need to make occasional adjustments using the slow-motion push buttons on the drive control, but these are much smaller and less frequent than with an alt-azimuth or an undriven equatorial. It is also worth checking the orientation of the image at the end of the observation as well as at the beginning, just in case you have accidentally jogged the projection screen, as sometimes happens.

When you have drawn all you can see on the full-disc solar image, put a plain piece of Bristol Board, cartridge paper or other high-quality paper over the grid and insert a higher-power eyepiece into your telescope, taking care not to disturb the orientation of your projection box or screen. For example, you might use a 9 mm Orthoscopic, giving a magni fi cation of 101×, for high-power work on the Sun with a 80 mm refractor.

Now carefully scan the enlarged solar image and look for any small spots you may have missed on the whole-disc image. When you see a hitherto unrecorded spot, make a mental note of its position relative to other sunspots, switch back to lower power and record its position on the full-disc drawing. More often than not, such missed sunspots are large enough to be easily visible on the smaller image once you know where to look for them. Sometimes a spot is big enough to make you wonder how you missed it in the fi rst place! When scanning at high power, pay particular attention to the eastern and western limbs of the Sun, where it is easy to miss small spots coming onto or moving off the Earthward side of the Sun. Often an area of faculae seen near the limb at lower power turns out to have one or more spots embedded within it when examined with higher magni fi cation. Also deserv-ing close scrutiny are the areas between the main sunspot latitude bands and the poles, where high-latitude sunspots occasionally appear, especially around the minimum of the sunspot cycle.

When you have fi nished scanning the image at high power and recorded any missed spots, the drawing is complete. The fi nal stage is to record some essential details: the date on which you made the drawing, the time at which you completed it, the observing conditions and the instrument used. Astronomers write the date in successively smaller units, e.g., 2012 July 10. The time should always be in Universal Time (UT), which is the same as Greenwich Mean Time. If Daylight Saving Time (Summer Time) is in use or if you live in a different time zone from GMT, remember to subtract the appropriate number of hours from the time shown on your watch. Some observers record the conditions in words – e.g., “clear,” “steady,” “boiling,” “passing clouds,” etc. – while others grade the steadiness of the seeing on a scale of 1–5, similar to the Antoniadi seeing scale used by planetary observers. In this scale, 1 is perfect or near-perfect seeing with hardly a ripple in the image (very rare from most sites), while 5 is extremely poor seeing, with the image “boiling” violently and making a drawing hardly worthwhile.

Some solar observing organizations also ask for observers to include the current rotation number with their drawings. This is a system of marking the Sun’s rotation and was begun in the nineteenth century by Richard Carrington (1826–1875). (Carrington was also co-discoverer of the fi rst solar fl are – see Chap. 3 .) Carrington’s system uses a rotation period of 25.38 days, with rotation number 1 beginning in

71Drawing Using the Projection Method

1853. Now, in the early twenty- fi rst century, the Sun has passed the 2,100th rotation since the commencement of Carrington’s system. The main use of the Carrington system is to determine the heliographic (i.e., solar) longitude of sunspots, as we shall discuss below.

Deriving Sunspot Positions

A completed solar drawing – or a correctly oriented photograph – can be used to derive the heliographic , or solar, latitudes and longitudes of sunspots and groups. From these data we can learn certain characteristics of the sunspot cycle and of the sunspots themselves.

Sunspot positions can be derived from drawings in two ways – by calculation or by using specially designed measuring grids. The mathematical method involves measuring the horizontal and vertical positions of the spots on your two-dimensional drawing or image and then translating these x - y coordinates into heliographic longitude and latitude on the three-dimensional solar disc using trigo-nometry. In the past, this method required a considerable amount of calculation and was very tedious, but now there are computer programs that do the mathematics for you – see, for example, the Helio Viewer software developed by British solar observer Peter Meadows and available for download from his excellent website, http://www.petermeadows.com/html/software.html . In this type of program, you can enter the x – y positions of the sunspots from your drawing or photograph and the software will calculate their heliographic coordinates.

A second, more traditional method of determining sunspot positions uses a series of measuring grids known as Stonyhurst discs , after the nineteenth- and early twentieth-century solar observatory at Stonyhurst, Lancashire, England, where they were invented. Stonyhurst discs consist of 6-in. circles with the cardinal points marked, as on our solar drawing grid described above, but instead of horizontal and vertical lines they are ruled with lines of latitude and longitude (Fig. 4.4 ). The verti-cal lines (longitude) become more curved with distance from the center of the disc, reminding us that the Sun is a three-dimensional object. To measure sunspot posi-tions, a completed solar drawing is laid over the Stonyhurst disc and the latitudes and longitudes of the spots are simply read off.

Both the mathematical and the Stonyhurst disc methods are as accurate as each other: whichever method you use, the data you put in derives from the original drawing, and so what matters is the accuracy of that drawing. A good 6 in. disc drawing can, in theory at least, give results accurate to 1° for spots near the center of the disc. Results are less accurate, though, for spots close to the limb. If you want to measure sunspot positions from your images rather than drawings, then you are probably better off downloading the appropriate software and using the fi rst method, as Stonyhurst discs are designed to be used underneath drawings made on thin paper, rather than with images printed on thick paper or viewed on a computer screen. But even if you only use images, it is worth learning how to use the

72 4 Solar Position Measurements

Stonyhurst disc method, as this gives you an appreciation of how heliographic coordinates are applied. Also, as described in Chap. 5 , Stonyhurst discs can some-times prove essential in visual sunspot counting, when you need to accurately determine the distances between groups. You can obtain a set of Stonyhurst discs from the solar observing section of a national astronomical society, or can down-load them from online sources – see, for example, www.petermeadows.com .

Determining heliographic coordinates is made complicated not only because we are translating two-dimensional coordinates into positions on a three-dimensional sphere, but also because the orientation of the solar image that we so carefully estab-lished above when making the drawing is only the orientation as seen from Earth . Earth’s axis is not perpendicular to Earth’s orbit but is tilted 23.5° to the perpendicu-lar. This, of course, is the cause of the seasons, but it also causes the apparent position of the Sun’s poles to vary. Similarly, the Sun’s own axis of rotation is inclined at 7.25° to the perpendicular to Earth’s orbit. As seen from Earth, the axial tilts of the two bodies combine to produce two effects. First, during the year the position of the Sun’s north pole varies by 26.3° either side of the apparent north point of the disc. Secondly, the Sun appears to “nod” towards and away from us by 7.2° (Fig. 4.5 ). We must cor-rect for both if we are to establish accurate sunspot positions.

Fortunately, both effects are regular and can be predicted accurately in advance. The changing angle of the north solar pole to the apparent north point of the solar disc is known as the position angle and is usually abbreviated to P . When the north pole is on the east side of the apparent north point it has a positive value and when on the west side it is negative. In early January the Sun’s axis coincides with the

Bo = 6.0°

-40°

-30°

-20°

-10°

0°

+10°

+20°

+30°

+40°

Bo = -6.0°

-40°

-30°

-20°

-10°

0°

+10°

+20°

+30°

+40°

Fig. 4.4 An example of a Stonyhurst disc, for a solar tilt ( B 0 ) of 6° (Courtesy of Peter Meadows)

73Drawing Using the Projection Method

apparent north-south line, and so the value of P is 0°. The axis then gradually moves west of apparent north, reaching its maximum westward tilt of 26.3° (i.e., −26.3°) in early April. After April the axis backtracks, returning to the 0° value in early July. It then tilts increasingly eastwards, reaching its maximum eastward value of +26.3° around October 11, before moving westwards again and returning to the zero point in January. Put more simply, the value of P is negative (i.e., west) in the fi rst half of the calendar year and positive (east) in the second half.

The extent of the Sun’s “nodding” towards or away from us is de fi ned as the heliographic latitude of the apparent center of the disc and is abbreviated B

0 (pro-

nounced B-nought). The value of B 0 does not vary in synchronization with P . In

early December the solar equator coincides with the center of the disc and so B 0 is

0°. In the early months of the new year the south solar pole tilts towards us, giving B

0 a negative value. The south pole reaches its maximum tilt towards us of 7.2° in

March and then begins to tilt away from us, passing through the zero point again in early June. For most of the latter half of the year the north pole is tilted towards us and B

0 has a positive value, reaching its maximum of +7.2° in September.

Thus, at most times of the year the lines of solar latitude – and the apparent paths followed by sunspots as they rotate across the disc – are curved, because the poles are at an angle to us. Only during a few brief days in June and December, when B

0

N

W

S

E

Early JanuaryP=0°; B0=−3.5°

P=0°; B0=+3.5°

P=−26.3°; B0=−6.2°

P=+26.3°; B0=+6.2°

N

W

S

E

Early April

N

W

S

E

Early July

N

W

S

E

Early October

Fig. 4.5 Illustrating the Sun’s changing position angle ( P ) and “nodding” towards and away from us ( B

0 ) during the course of the year (Courtesy of Dominic Ford)

74 4 Solar Position Measurements

is zero, do sunspots follow straight paths across the disc. Therefore solar observers use a set of eight Stonyhurst discs, one for each degree of B

0 . Note also that P and

B 0 have one feature in common: the rate of change of the value of both angles is

slowest when they are at their maximum and fastest when they pass through the zero point. For example, P is more than −26° for over 3 weeks from late March until mid-April, whereas in January it varies by nearly 5° (passing through 0° on January 5) during the fi rst 10 days of the month.

A knowledge of P and B 0 at the date of observation is suf fi cient for calculating

the latitudes of any spots present. However, to calculate the longitudes of spots we need a third statistic. Because the Sun’s surface is gaseous and there is no fi xed point from which to reckon longitude, solar astronomers measure longitudes from the central meridian . This is simply a line drawn across the disc from the position of the north end of the rotation axis (as determined by the value of P ) through the center of the disc and is the point at which sunspots and other solar features pass across, or “transit,” the disc’s center. The longitude of the central meridian (abbrevi-ated L

0 ) is 0° at the beginning of each new solar rotation number, as determined by

Carrington’s system described above. Heliographic longitude is measured from east to west, but the Sun also rotates from east to west, and so the longitude decreases with time, progressing from 0° to 350°, then to 340° and so on.

Example

To demonstrate how to work out P , B 0 and L

0 , and thus derive the heliographic

latitudes and longitudes of sunspots using Stonyhurst discs, let us use one of this author’s own drawings, made at 13:28 UT on October 12, 2008 (Fig. 4.6 ) as an example. A drawing made at sunspot minimum was deliberately chosen, in order to keep things simple. At maximum, when there are often many large, complex groups on the disc, measuring the positions of all the sunspots would take much longer.

To begin with, we need to fi nd the three essential parameters for the date and time at which the drawing was made: P , B

0 and L

0 . These are tabulated in several

annual publications, including the Handbook of the British Astronomical Association and the Astronomical Almanac . Both publications are released annually. You can also obtain current values of P , B

0 and L

0 from some online sources – for example,

http://www.jgiesen.de/sunrot/index.html – though web resources can be more ephemeral than the major annual astronomical publications. The Astronomical Almanac lists the fi gures for each day of the year, while the BAA Handbook gives them at 5-day intervals, plotted for the previous midnight, (i.e., 0 h UT). P and B

0

change slowly enough that we can use just the daily fi gure, and so users of the Almanac can simply copy down the values for the date of the drawing. If you use the BAA Handbook you need to interpolate between the two nearest values. Because the BAA Handbook is the most commonly used reference source in the UK, where this author lives, it will be used in our example.

75Drawing Using the Projection Method

October 12, 2008, the date of our example drawing, falls between two of the BAA Handbook 5-day intervals: October 11 ( P = +26.3°) and October 16 ( P = +26.1°). To fi nd the value of P for October 12, we fi rst need to fi nd the differ-ence between the values for October 11 and 16 (26.3° − 26.1° = 0.2°) and divide the result by 5 (giving 0.04°). This means that P is changing by increments of 0.04° per day between the two dates. As P on October 16 is less than on October 11, P is decreasing, so we need to subtract from the value for October 11. October 12 is just 1 day after this, so we need to subtract just one of the fi ve daily increments of 0.04° from the value for October 11 – i.e., 26.3° − 0.04° = 26.26°, which can be rounded back up to +26.3° if we are calculating to just one decimal place. Note that +26.3° is as large as P ever gets, the Sun having reached its maximum eastward tilt only a few days earlier in October.

To determine B 0 we need to repeat this process. B

0 for October 11, 2008, was

+6.1°; on October 16 it was +5.8°. 6.1° − 5.8° = 0.3°. 0.3° divided by 5 = 0.06°. Therefore B

0 for October 12th = 6.1° − 0.06° = 6.04°, which can be rounded down

Fig. 4.6 Whole-disc drawing, made by the author on October 12, 2008, showing a small bipo-lar sunspot group. The position angle P is marked

76 4 Solar Position Measurements

to +6.0°. Be very careful to get the signs right when calculating P and B 0 , as

otherwise the results will be hopelessly wrong – although, as discussed above, you can generally reckon on P and B

0 being negative in the fi rst half of the year and

positive for most of the second half. Because the Sun rotates through a full 360° in 25.38 days L

0 , the longitude of

the central meridian, changes quickly enough that its value for the time of day on which you make your observation is signi fi cantly different from the midnight value. L

0 decreases by 13.2° per day, which means that it decreases by 0.55° per hour. To

obtain the value for 00:00 UT on October 12 you need to subtract 13.2° from the tabulated value for 00:00 UT on October 11 – i.e., 171.3° − 13.2° = 158.1°. The “Sun” section of the BAA Handbook includes a table headed “Decrease of L

0 with

time.” From this you can see that during the fi rst 12 h of the day the value has dropped by 6.6°. To get the amount of decrease for 13:28 UT on October 12, you need to subtract this 6.6° from the midnight value, and then subtract a further incre-ment representing the decrease from 12:00 to 13:28 UT, i.e. for 1 h 28 min, which from the table is 0.8°. Therefore the total decrease in L

0 since 00:00 UT on October

12 is 158.1° − 6.6° − 0.8° = 150.7°. If you use the Astronomical Almanac , you need to interpolate between the values

for midnight on October 12 and the same time on October 13 to obtain L 0 for the

time of day at which you are observing. Because you need to refer to them, it is useful to make a note of P , B

0 and L

0 on the sheet on which the drawing is made.

The next job is to mark the Sun’s axis on the drawing. We have found the value of P to be +26.3° – that is, the north end of the Sun’s axis is inclined 26.3° east of the apparent north point of the disc. Mark the appropriate angle on the drawing using the 6-in. protractor and draw an accurate line across the solar disc through this mark and the cross at the center of the disc. In the example given in Fig. 4.6 this has already been done. A 6-in. protractor will not allow you to measure to one-tenth of a degree, so just mark the angle to the nearest half-degree.

We are now ready to use the Stonyhurst discs. Select the disc whose B 0 value is

the closest to the B 0 value for the day of your observation. In this case, B

0 = +6.0°,

so you should choose the disc for 6°. As B 0 is positive, the northern hemisphere

(sometimes marked by a + sign before the value of B 0 ; otherwise the one without a

− sign) should be at the top. Lay your drawing over the disc, turning it around so that the Sun’s axis which you marked on the drawing coincides with the central, straight north-south line on the disc (Fig. 4.7 ). Make sure that your drawing disc coincides exactly with the edges of the Stonyhurst disc, and then clamp the two sheets fi rmly together with a paperclip on at least two sides.

It is very satisfying to look at your drawing at this point, because the Stonyhurst disc’s curved lines of latitude and longitude give it a three-dimensional appearance, reminding us that we are drawing and measuring positions on a sphere. Because the line of the equator is accurately established, we can tell instantly in which hemisphere a spot is located. On this particular day one small sunspot group was visible, located in the northern hemisphere – proof that even at solar minimum there are often inter-esting things to see on the Sun – though these were the fi rst sunspots that had been seen for several months during the unusually deep sunspot minimum of 2008.

77Drawing Using the Projection Method

Reading off the latitude of the group is now a simple matter, because on the Stonyhurst disc latitude is marked off at 10° intervals and (in some versions of the discs) for every degree at the limb. Measure the latitude of a spot using the center of its umbra. You can use a ruler graduated in millimeters to measure the latitudes more accurately. For example, the leader sunspot of the group in the drawing is 7.5 mm above the line representing +20° latitude. The distance between the 20th and 30th latitude parallels is 13 mm. Therefore the spot is (7.5/13 mm) × 10° = 5.8° above the 20th parallel, so its latitude is +25.8° or 25.8°N. Measuring in the same way, we fi nd that the follower spot of the same group has a latitude of +28.1°.

To measure the longitude of a spot we need to know its distance from the central meridian – i.e., the north-south axis line. On the Stonyhurst discs the longitudes are marked off at intervals of 10°, as for latitude. It is usually quite easy to interpolate longitudes to the nearest degree using a millimeter ruler, as with latitude. The essential thing to remember when measuring longitudes is that they run from east to west , so longitudes of spots west of the central meridian are higher than that of the central meridian itself (i.e., L

0 ), whereas those for spots east of this line are

Fig. 4.7 The author’s drawing of October 12, 2008, shown laid over a Stonyhurst disc repre-senting B

0 = 6.0° (Stonyhurst disc courtesy of Peter Meadows)

78 4 Solar Position Measurements

lower. Thus, in our example drawing, to fi nd the longitude of the leader spot of the group, which is east of the central meridian, we need to measure its distance from the line (28.2°) and subtract this value from L

0 (150.7°), giving a heliographic

longitude of 122.5°.

What We Can Learn from Drawings

We noted in Chap. 3 that sunspots form in distinct “bands” of latitude on either side of the equator – generally between 5° and 30°. Spots at latitudes of higher than 40° are quite uncommon and never form at or near the poles. But these bands are not constant. If we make measurements of sunspot positions over several years and plot the latitudes of the spots on a graph, we fi nd that the average latitude at which spots form gradually decreases as the solar cycle progresses. Early on in a new solar cycle, shortly after minimum, spots tend to appear at latitudes well above 20° and often at 30° or even higher.

By the time of sunspot maximum the average latitude of sunspots has declined to about 15°, and as we approach minimum the small number of spots seen at this stage of the solar cycle usually form within a few degrees of the equator. However, by the time activity “bottoms out” at minimum, occasional spots are seen at high latitudes again, heralding the start of the new solar cycle. In this respect there is always a certain amount of overlap between one cycle and the next. This phenom-enon of gradual decline in sunspot latitudes was discovered in the nineteenth century by (yet again) Richard Carrington. Further studies into it were undertaken by Gustav Spörer (1822–1895) in Germany, and today it is known as Spörer’s law – rather unfairly, as there is no doubt that Carrington discovered it fi rst.

The work of Carrington and Spörer was extended in the early 1900s by Edward Walter Maunder (1851–1928), who was for many years head of the Solar Department at the Royal Observatory in Greenwich, England. Maunder plotted the decline of sunspot latitudes over a period of many years, covering several sunspot cycles, and the resulting graph is known as the “Butter fl y diagram” on account of the shape generated by each sunspot cycle (Fig. 4.8 ).

Our example drawing used above to describe how to measure sunspot positions provides a good demonstration of Spörer’s law. This drawing was made around the time of minimum activity in 2008, and the leader spot’s latitude of +25.8° is a high latitude, indicating that the group belongs to the new cycle that was just starting. It was, in fact, the fi rst sunspot group of the new cycle that I saw – though the fi rst new-cycle spot seen by anyone appeared in January of that year. Conversely, we had not seen the last of the low-latitude spots of the old cycle by that October day in 2008; indeed, a few such spots were visible as late as August 2010.

Solar minimum can be a trying time for the solar observer, since there are some-times no sunspots for days or weeks on end, and, with the exception of the polar faculae, there is often very little to look at on the Sun, at least in ordinary white light. It is therefore rather exciting to see and record the fi rst high-latitude sunspots,

79Drawing Using the Projection Method

as this tells us that the next solar cycle, and therefore greater activity, is just around the corner – although activity was very slow to pick up after the 2008 minimum, and it was only during the 2010–2011 period that the new cycle really took off.

For convenience, astronomers refer to individual sunspot cycles using a sequen-tial series of numbers, starting at solar cycle 1, which reached maximum in 1761. Choosing this cycle as number 1 was an arbitrary decision by astronomers, as of course activity was varying in cycles before then. Now (2011) we are about half-way to the maximum of solar cycle 24. The high-latitude spot shown in our example drawing was one of the fi rst spots of cycle 24, which began in 2008 and is expected to reach maximum around 2013. Minimum, and the start of cycle 25, are expected to occur around 2019. These forecasts assume that cycle 24 will have an average length. We have already seen that the cycle got off to a slow start, and only time and patient recording will tell whether this tardiness will be a characteristic of the new cycle or whether the rest of the cycle will show more conventional behavior.

Measurements of the longitudes of sunspots can also be useful. For example, as noted in Chap. 3 , some large and active sunspot groups survive more than one passage across the Sun’s visible hemisphere. If we know such a group’s latitude and longitude on its fi rst rotation, we can very quickly establish whether or not it is the same group when it reappears at the eastern limb on its second passage. Although we can now track active regions all round the Sun using spacecraft images online, it is nevertheless very rewarding to prove using your own drawings that a group is visible the second time around. Sometimes the second appearance of a group is obvious from its general shape and approximate position on the Sun’s

Fig. 4.8 The “Butter fl y Diagram,” plotted from the 1880s up to 2011 ( top ), matched up with the graph of sunspot numbers for the same time period ( below ). The migration of sunspots from high to low latitudes in each cycle is clearly seen (Courtesy of NASA/David Hathaway)

80 4 Solar Position Measurements

limb, plus the knowledge that it disappeared around the western limb 2 weeks before, but often it takes a precise plot on both the fi rst and second passages to pin it down for sure.

Another useful project is to plot the longitudes of the spots against time on a graph, as for the latitudes. You can do this for just 1 month or one solar rotation or for several years. Using this technique, observers have found that some longitudes are sometimes more active than others.

Detailed Drawings

If you have time, it can be worth making a more detailed drawing of an interesting spot or group at higher magni fi cation. You can make such drawings using a grid as you would for a whole-disc plot, but with a blank sheet of paper placed over the grid on the clipboard. You could prepare a small grid specially designed for draw-ing individual sunspot groups, but it is actually just as easy to project the high-power image onto the grid used for the whole-disc drawings and use that to draw the spots relative to each other. If time is short, you can just make a freehand sketch. The essential thing to remember when making detailed drawings is to start by draw-ing the positions of the major sunspots on the grid squares fi rst and add the fi ner details later.

Do not try to draw in too much detail fi rst, as the result will be less accurate. A well-aligned equatorial mount with an ef fi cient motor drive is especially useful for making high-power drawings. Remember to accurately record the directions of north and east, so as to establish the orientation of the drawing, and also the same details as you would for other types of observations – date, time, seeing, instru-ment, magni fi cation, etc.

For most purposes, detailed drawings such as this have been superseded by electronic imaging. An image can record the positions of sunspots relative to each other (as opposed to their positions relative to the points of the compass) more accurately than drawing, and the superior resolution of good images taken with digital cameras and webcams obviates any advantage in resolution that drawings ever had over traditional photography.

However, for many of us, a detailed sketch offers a cheap and convenient way of recording sunspot detail. In any case, it is a good idea to know how to do it, as sometimes it is the only way in which to record an interesting feature before clouds roll in or a very rare phenomenon such as a white light fl are appears. White-light fl ares last for only a few minutes and could easily fade from view before you have set up your camera and fi lter. On the other hand, you can begin drawing such a feature pretty much as soon as you see it. Sketching is also a good way of getting a quick record of transient features in H-alpha when you don’t have time to set up imaging equipment (see Chap. 6 ).

81Cooperation with Other Observers

Cooperation with Other Observers

In an ideal world the amateur astronomer would observe the Sun every day, as only then can a complete picture of the Sun’s behavior be obtained. But in reality no one person can monitor the Sun every day. Depending on your local climate, some days will be cloudy, while the restrictions imposed by our daily lives can severely limit the time available for solar observation. In particular, if you live in a high temperate latitude such as the UK or parts of the United States and Canada, you may well fi nd that you cannot observe the Sun at all during weekdays in winter. You might regu-larly manage to observe the Sun on 15–20 days during the summer months, because the Sun is still up when you get home on weekday evenings, but in winter you may do well to make observations on more than 5 days per month.

In any case, the Sun is by de fi nition only visible during daylight, and so a con-stant watch on the Sun can only be maintained by observers scattered around the globe. You can help in the effort to monitor the Sun by sending your observations in to a central coordinating body, which pools together results obtained by a large number of observers and produces reports on the Sun’s activity. Observing as part of a team applies not only to drawings and sunspot positions but also to sunspot counts and other observations described elsewhere in this book. Coordinating bod-ies for observations are often the solar sections of national amateur astronomy organizations, and several such groups are listed in Appendix C of this book. You might want to send your observations to the solar sections of the Society for Popular Astronomy and the British Astronomical Association if you live in the UK. The latter has over 50 regular observers operating from many countries, and for many years the section has managed at least some coverage of the Sun on every single day of the calendar year.

83L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0_5, © Springer Science+Business Media New York 2012

Chapter 5

The most scienti fi cally useful, and in some ways the easiest, type of solar observation an amateur astronomer can make is to measure the level of activity seen on the Sun. This is best done by counting the number of sunspots and/or sunspot groups visible on the Sun’s disc each day. Other methods of measuring solar activity are used by professional observatories, but sunspot counting is the most practical option for the amateur.

Sunspot counts are very useful in monitoring the Sun’s behavior, and one type of sunspot count, known as the Relative Sunspot Number, is an internationally recog-nized measurement of sunspot activity. Measurements of the Relative Sunspot Number go back well over a century, and they have been reconstructed for even earlier times by astronomers going through old observations, so when you have learned the correct methods of sunspot counting you can contribute to a long histori-cal record that is accurate and consistent and allows us to compare what the Sun is doing now with what it was doing, say, 100 years ago, long before the sophisticated methods of monitoring the Sun used by today’s professional observatories were available. Although most scientists do not believe that the Sun is the main cause of contemporary global warming (see Chap. 1 ), long-term studies of solar activity such as sunspot numbers are obviously invaluable in research into climate change.

Sunspot counting is quite easy, and a simple daily sunspot count can take just a few minutes. This is good news if you are pressed for time; although drawing the Sun is very satisfying to do, it can be quite time-consuming, especially if the Sun is active. For the same reason, sunspot counting is very suitable if you live in a cloudy climate, as in such conditions the Sun is often visible for only a short clear spell between clouds. It can be very frustrating to have to abort a drawing or imaging session half-way through because of clouds. For this reason, it is recommended that you do sunspot counting fi rst during a solar observing session, because if clouds

Measuring Solar Activity

84 5 Measuring Solar Activity

eventually roll in and curtail drawings, photographs and so on then you at least have something to show for your day’s solar observing.

Another major advantage of sunspot counting is that you can do it equally well whether you project the Sun’s image or view it directly using a fi lter, since it does not require plotting or accurate position measurements. It is therefore equally suitable for observers with Schmidt-Cassegrain or Maksutov telescopes as it is for those using traditional refractors and re fl ectors.

Try to make sunspot counts on as many clear days as you can. This is partly so that you can gain practice at the various methods of counting, but also because the more data you gather, the more accurate a picture of the ebb and fl ow of the sunspot numbers you will build up. As with drawings, assembling a complete picture of the Sun’s behavior is only possible if you make observations in collaboration with other astronomers, and we will describe methods of reporting results to solar observing organizations below. Maintaining daily observations is especially important in periods of bad weather, as only a small number of observers may be active then, and your results could be important.

The Mean Daily Frequency

The simplest form of sunspot counting is to count the total number of sunspot groups – sometimes known as “active areas” – seen on the Sun’s disc each day. For example, the image of the Sun shown in Fig. 5.1 reveals three groups, so the count of active areas for that day should be recorded as 3. At the end of the month, add up the numbers of groups and then divide the total by the number of days on which you made sunspot counts. The resulting average is known as the Mean Daily Frequency (MDF) of active areas.

To understand how the MDF works let us use an example from the author’s own observing notebook. Table 5.1 lists counts of active areas made in July 2005. The left-hand column gives the day of the month on which an observation was made, and on the right is the number of active areas seen on that day.

Adding up the fi gures on the right-hand column gives us 39. Counts were made on 13 days in July 2005, so if we divide 39 by 13 we get an MDF of 3.0. This is a quite mediocre MDF, indicating only low to moderate solar activity – just what we would expect in an observation made some 5 years after the last solar maximum, when the solar cycle was well on its way towards minimum. Note, though, that even at this stage of the cycle activity can be highly variable. As many as seven active areas were recorded on 2 days at the beginning of the month, whereas later in the same month the sunspot count was zero. Note also that if you see no spots you must always record it as zero and not simply leave the record blank for that date, as a zero count is as important as any other fi gure and must be included in the calcula-tion if the MDF is to be accurate.

Of fi cially, any spot or group of spots that is at least 10° of heliographic latitude or longitude away from any other spot counts as one active area. The size or com-plexity of a group does not matter. Whether it is a single, small spot or a large,

85The Mean Daily Frequency

Fig. 5.1 Whole-disc image of the Sun, showing three sunspot groups. The active area count for that day should therefore be recorded as 3 (Photograph by the author)

Table 5.1 Counts of active areas made by the author in July 2005, showing how to work out the Mean Daily Frequency (MDF)

Day Active areas

2 7 5 7 8 4 10 4 11 3 12 5 14 2 15 2 17 0 18 0 20 0 21 1 28 4 Total 39 No. of observations 13 MDF 3.0

86 5 Measuring Solar Activity

complex group, it counts as one active area. For a group to be counted, however, it must contain at least one sunspot with a dark umbra. An isolated area of faculae that does not contain any sunspots is not included in an active area count. Neither are pores (described in Chap. 3 ), isolated pieces of penumbra or the occasional grayish patches included.

Two exceptions to the above “10° rule” occasionally occur. One is when a large group expands to the point where it is more than 10° long and contains two distinct centers of activity more than 10° apart. Such a group should then be counted as two active areas. The 10° rule may make it seem dif fi cult to determine what counts as an active area, but in practice the vast majority of groups are very distinct and sepa-rated by well over 10°. If a borderline case does occur, however, the best way of ascertaining whether it is one or two active areas is by making a disc drawing (as described in Chap. 4 ) or taking a whole-disc picture. Then, using the appropriate Stonyhurst disc, it is easy to measure the separation between the groups and deter-mine whether they are more than 10° apart. For this purpose, you can sometimes insert a Stonyhurst disc into your projection box and measure the distance between two groups directly off the projected image. If you don’t have time to make a draw-ing or a measurement you must use your own judgment, but in any case you should always make a written note of your decision.

When making an active area count, begin by counting all the groups you can see on the Sun’s disc, using a low magni fi cation. Then, switch to a high-power eyepiece and carefully scan the Sun in north–south or east–west strips. As noted in Chap. 4 on making solar drawings, you will often be surprised at how many spots you pick up that were not visible at low power. Be especially vigilant when searching the eastern and western limbs of the Sun, and also when scanning the high-latitude zones around 40°, where spots are relatively rare but small ones occasionally appear, especially around sunspot minimum. Because every isolated spot, however small, counts as an active area, a large number of small spots, unde-tected until you have scanned the Sun at high power, can make a big difference to your active area count.

You can, in fact, make a reliable active area count from a drawing, at your leisure after the observation, provided you have carefully examined the Sun at high power and recorded the positions of all the spots present. As remarked above, however, it is often not safe to do it this way in many climates, where clouds can roll in before you have completed your observation!

The Relative Sunspot Number

The standard method of counting sunspots used by astronomers throughout the world is somewhat more sophisticated. Known as the Relative Sunspot Number , it takes into account both the number of sunspot groups (active areas) present, counted using the method described above, and the number of individual sunspots within these groups. The Relative Sunspot Number was started way back in 1848

87The Relative Sunspot Number

by J Rudolf Wolf (1816–1893), later director of the Swiss Federal Observatory (now the Institute for Astronomy) in Zurich, and has been in use ever since. It is sometimes known as the Wolf Number or Zurich Number . Zurich became an inter-national center for determining the daily Relative Sunspot Number, coordinating sunspot data contributed by amateur and professional astronomers worldwide.

The Institute for Astronomy relinquished this responsibility in 1980, and nowa-days the central coordinating body for the sunspot number is the Solar Influences Data Analysis Center in Brussels (part of the Royal Observatory of Belgium), where an of fi cial sunspot number for each day, known as the International Sunspot Number , is worked out. A similar system operates in the United States to produce an American Relative Sunspot Number. Amateur solar observing groups worldwide send their results to both organizations. It gives the aspiring solar observer some encouragement to know that Wolf made his original observations with an 80 mm f/14 refractor – exactly the sort of instrument employed by modern amateur solar observers – and, even more remarkably, that the same telescope is still used daily to count the sunspot number! One of the great strengths of the Relative Sunspot Number is that it continues an unbroken record of solar observations, made with the same type of instrument, going back over 150 years.

The daily Relative Sunspot Number, usually abbreviated to R , is worked out using the following formula:

( )= +10R k g f

where g is the total number of sunspot groups (active areas), counted in exactly the same way as for the MDF, and f is the total number of individual sunspots within the groups. The g is multiplied by 10 in the formula because Wolf noted that there were approxi-mately ten times as many individual spots as there were groups, and so multiplied the statistical weighting of groups by ten to give them fair representation. The k is a constant included to standardize results from a wide variety of observations.

Sunspot counts vary widely between observers for many reasons, including the size, type and quality of the telescope, local seeing conditions and the ability and experience of the observer. Determining your own constant requires making sun-spot counts over a long period and then comparing your results to the of fi cial, published values for the same period. However, solar observing organizations often do this for you, and so when calculating your own Relative Sunspot Number you should not include a constant unless instructed to do so, and simply use the formula = +10R g f .

Let us use a few examples to show how R is determined. Figure 5.1 shows three active areas on the Sun. The group on the right contains four spots, the group at the center with the large leader spot has fi ve spots, while the group at left contains just two spots. The total number of individual spots f is therefore 4 + 5 + 2 = 11. We need to add this fi gure to the total number of groups g multiplied by 10, i.e. 3 × 10 = 30. Therefore our fi nal R is 30 + 11 = 41. The smallest possible value for R that is not zero is when there is just one single spot on the entire solar disc. R is then 11 – i.e., 1 group + 1 spot.

88 5 Measuring Solar Activity

Sometimes you will see a published value for R that is less than 11. This can happen because when there are just one or two tiny spots, not all observers will see them – some observers have better seeing conditions than others, some observers might not notice the spots, and the spots may be short-lived and so are simply not present when some contributors are observing – and so the value for R averaged for all observers will be reduced. Introducing a constant can also sometimes reduce the value. This illustrates the importance of using a large number of observers in order to build up the most accurate measurement possible of the Sun’s activity. At the other end of the activity scale, one of the highest levels of activity this author has ever seen was on July 19, 2000, when there were no less than 16 groups and 193 individual spots, giving an R of 353.

To do the sunspot counts to enable you to calculate R , begin by counting the number of active areas exactly as you would for an MDF count, remembering to scan the Sun at high power for small spots. Then, again using the high-power eye-piece, count all the spots that you can see within the groups. Only at a relatively high magni fi cation are individual spots within groups satisfactorily resolved. For your sunspot counts, you can use two eyepieces as you might use for drawings: for example, a 15 mm (61×) for an initial reconnaissance of the whole disc and a 9 mm (101×) for a detailed scan of the Sun for small spots and counts of individual spots within the groups. Always count the groups fi rst and then the spots, for the same reason as it is advisable to do sunspot counts before drawings or imaging: if during the spot count it clouds over for the rest of the day, you have at least got a basic MDF count.

To count the number of spots within a group correctly, it is important to know what is of fi cially a spot and what is not. To be included in a count a spot must contain an umbra and must be entirely separate from any other umbrae in a group (Fig. 5.2 ). Two or more umbrae within the same penumbra should be counted as two or more spots – as long as the umbrae are completely separated from each other. If you see what looks like two spots joined together, always count them as just one spot unless they are entirely separated by bright photosphere between them. If a light bridge appears in a spot and divides it into two, then you should count it as two spots – but only when the division is complete. A spot with a light bridge running only part of the way across it is still one spot – although you should make a note that a light bridge is developing.

As with MDF observations, you should not include pores and penumbrae in spot counts. When a group is near the eastern or western limb, or when seeing condi-tions are less than good, it is often dif fi cult to distinguish pores and small sections of penumbrae from genuine umbrae. Small pieces of penumbra are often found around the edges of large, complex groups, particularly when they are starting to decay. When deciding whether a feature is a genuine spot, compare how dark it appears with nearby “real” sunspots. A genuine umbra appears very dark, while penumbrae are a lighter gray when examined closely.

At times of high activity, when there are groups containing large numbers of individual spots, f (the total number of spots on the disc) can sometimes be 100 or more, and it is easy to lose count when tallying them up. One way to get around

89The Relative Sunspot Number

this is to note down the number of spots in each group and then add up the total at the end of the observation. In any case, it is often a good idea to count f twice, if possible, in order to check that you have not made any major errors in counting.

When you have completed your sunspot count, do not forget to record the time at which you fi nished the observation (in Universal Time [UT], to the nearest min-ute) and the seeing conditions, perhaps using the same 1–5 seeing scale described in Chap. 4 . It is also important to note down full details of the telescope used, including whether you used projection or observed the Sun directly with a fi lter. If you observe the Sun directly, you should record the magni fi cations of the eyepieces used, while if you use the projection method you should also note down the approx-imate diameter of the Sun’s disc as seen on the projected image.

You can get the disc diameter for a high-magni fi cation projected image by divid-ing the magni fi cation of the high-power eyepiece by that of the lower one, and then multiplying the result by the diameter of the whole-disc mag used for low-magni fi cation work. For example, if you use an eyepiece giving a magni fi cation of 61× to produce a 152 mm solar image, then the disc diameter produced by a 101× magni fi cation on the same telescope and assuming the same projection distance between eyepiece and screen is (101/61) × 152 = 251.7 mm. If you stick with the same instrument and projection setup, then you need only record your instrument details once, but always record any changes you make to your instrument.

Fig. 5.2 This image of a sunspot group gives a good demonstration of which internal spots should and should not be counted as spots for the purpose of determining the Relative Sunspot Number, R . All the dark umbrae towards the left and right ends of the group should be counted, but the pores and gray penumbral material at the center should not (Image by Dave Tyler)

90 5 Measuring Solar Activity

For a slightly more detailed report, you can divide your counts of both active areas and R into hemispheres. The northern and southern solar hemispheres often differ widely in the level of activity over the course of a month and sometimes over longer spells, and amateur solar observing groups publish MDF and R fi gures for each hemisphere as well as general data for the whole Sun.

Knowing what is the true northern and southern hemisphere on the Sun requires a knowledge of P (the angle of the Sun’s axis to true north) and B

0 (the Sun’s tilt

towards or away from Earth), as you would need when making position measure-ments from drawings or images. Although it is always best to make a drawing or take a photograph if you have the time, you do not always need to draw or image the disc to establish the positions of the hemispheres for sunspot counting purposes.

Because sunspots tend to occur in distinct bands of latitude in each hemisphere, it is often obvious which spots are in the northern hemisphere and which are south-ern, provided the image is correctly oriented as described in Chap. 4 and you know the value of P . Occasionally, however, spots do form close to the equator, at lati-tudes of 5° or even less. This often occurs around solar minimum and the years immediately preceding it when, in accordance with Spörer’s law (see Chap. 4 ), sunspots of the dying solar cycle form at lower latitudes. You may then need to make a drawing to establish which hemisphere such low-latitude spots are in, but it is usually possible to assign a spot to the correct hemisphere directly from the pro-jected image, with the help of Stonyhurst discs.

To establish the orientation of the Sun’s axis, look up the value of P in the BAA Handbook , the Astronomical Almanac or one of the online resources mentioned in Chap. 4 and Appendix D of this book. Once you have established the position of apparent north by letting a sunspot drift along the east–west line, imagine that the north–south line is tilted by the appropriate value of P . The east–west line is tilted by the same amount, and this roughly marks the division between the hemispheres (Fig. 5.3 ). To establish the angle with any degree of precision, you need to mark the circumference of the disc on your projection grid with angles at 5° intervals.

Unless you are observing during the late stages of a solar cycle, it should now be clear which hemisphere most of the spots are in. You may notice, however, that a few spots seem to pass through the exact center of the disc – an unlikely occur-rence, given that spots exactly on the equator are rare. This is due to the Sun’s tilt towards or away from us, as determined by the value of B

0 . If B

0 is positive, the

Sun’s north pole is tilted towards us, and so the equator curves to the south of the center of the disc (i.e., below it in most projected images), while if it is negative the equator curves to the north of the center. You need to take this into account when assessing which hemisphere a spot is located in, and the best way of doing this is to select the appropriate Stonyhurst disc for the current value of B

0 and place

it on top of your projection disc, tilting it over to the current value of P and remem-bering to orient the Stonyhurst disc appropriately according to whether B

0 is positive

or negative. It should then almost always be obvious whether a spot is in the north-ern or southern hemisphere. If you are still not sure, use your judgment to assess which is the most likely hemisphere and record in your report that the spot in question was so close the equator that you were not certain which hemisphere it was in.

91The Relative Sunspot Number

The business of assessing the Sun’s true orientation by eye may sound rather obscure, but it is actually quite easy after a little practice. However, if you are still not con fi dent about establishing north and south without making a drawing, you can leave this out of your sunspot counts, as the essential statistics are the total counts of active areas and spots on the whole solar disc.

It is important that you record your sunspot counts in a standard format, as this ensures that you record all the correct details. It also allows you to easily transfer your results onto the report forms provided by solar observing groups and enables you to compare your own observations directly with the results published by these groups.

The best way of recording sunspot counts is in a table for each month’s observa-tions, with a row for each day of the month, and Fig. 5.4 shows a sample page from the author’s own solar notebook. The standard table was prepared on a computer

TrueNorth(P)

Approx

TrueEquator

10°20°

30°

10°20°

30°

N

S

EW

Fig. 5.3 Diagram showing how to determine the Sun’s true orientation by eye. Note also that the Sun’s true equator is usually curved northwards or southwards, and this needs to be taken into account when determining which hemisphere a spot is in (Drawing by Dominic Ford)

92 5 Measuring Solar Activity

using Microsoft Word and a fresh pre-printed table is pasted into my notebook for each month, but you can draw one out by hand just as easily. It is a good idea to design your table to resemble as closely as possible the report forms distributed by the observing organizations to which you will send observations – such as the solar

Fig. 5.4 Sample page from the author’s solar notebook, showing sunspot counts recorded in tabular form

93The Relative Sunspot Number

sections of the British Astronomical Association and the Society for Popular Astronomy. The headings of the columns are as follows:

D Day of the month UT Universal Time at which the sunspot count was completed S Seeing (on the 1–5 scale described in Chap. 4 ) g

n Number of active areas seen in the northern solar hemisphere

f n Number of individual spots seen in the northern solar hemisphere

g s Number of active areas seen in the southern solar hemisphere

f s Number of individual spots seen in the southern solar hemisphere

g Total number of active areas seen (found by adding g n to g

s )

f Total number of individual spots seen (found by adding f n to f

s )

R Relative Sunspot Number, i.e., 10 g + f P The total number of prominences observed with an H-alpha fi lter (see Chap. 6 )

There are two additional rows at the bottom of the table. The fi rst is the total for each column, added up at the end of the month. This facilitates fi nding the average number of active areas seen (i.e., the MDF) and Relative Sunspot Number, which is written in the fi nal row. For quick reference you might also fi nd it helpful to make a note at the bottom of the page the month’s MDF, R and the number of days on which you made observations. (Here you might also record any prominences – counting prominences is described in Chap. 6 .).

A sunspot count table does not have to be as elaborate as this. You can divide your count up into hemispheres, which increases the number of columns, but if you wish to do sunspot counts for just the whole disc then you only need fi ve columns: two for g and R , in addition to those for the day of the month, the time and the see-ing conditions.

In addition to recording the fi gures, you might want to make a few brief notes on what you have seen each day, especially if you have seen something interesting or unusual – a new large group at the eastern limb, a spot split by a light bridge or an especially complex or active sunspot group, for example. Such notes are especially useful as an aid to memory if you did not make a drawing or were unable to observe the Sun for several days afterwards. They can also be important in their own right. For instance, at 12:45 UT on February 15, 1992, this author noticed a fairly small spot near the center of the Sun’s disc whose umbra appeared to be broken into sev-eral fragments. I did not make a drawing that day, but made a note of the spot’s unusual appearance and included it with my monthly report to the BAA Solar Section. Some weeks later, I read in the Solar Section’s monthly newsletter that several other observers had noticed unusual activity in the spot and that the follow-ing day the same spot had produced a white-light fl are – an extremely rare event, as described in Chap. 3 . If only I had observed the Sun the next day as well!

Nevertheless, this event shows how important it is to make notes, as you never know how useful they might be. More recently, in 2009, on two occasions during the summer of that year I recorded a very small, isolated spot, making a note of its position (and including it in the sunspot count, of course). On both occasions, the

94 5 Measuring Solar Activity

solar section director of the Society for Popular Astronomy reported back that I was the only section member to see each of these groups, though both were con fi rmed as active areas by the Mount Wilson Observatory website.

It is important to send your sunspot data to your solar observing organization promptly after the end of the month, so that it can collate and publish the results obtained by many observers as quickly as possible. The quickest way of reporting is, of course, by e-mail. You should use one of your group’s standard online report forms; if you use your own form, it should resemble the group’s form as closely as possible. Every month you can e-mail your completed sunspot count form to each organization, together with a separate Word document containing notes on your observations. You really don’t need to send drawings nowadays; instead, you can include with your observing notes latitude measurements derived from drawings.

Sunspot counts, especially of R , are the most important kind of solar observation amateur astronomers can make, and high-quality sunspot data has real scienti fi c value, continuing an unbroken record of more than 150 years of observations. The more observations received, the larger is the statistical sample and so the more accurate is our picture of the Sun’s behavior. The BAA and some other national solar observing organizations submit the best of the data received from observers to the Solar In fl uences Data Analysis Center (SIDC) in Brussels, Belgium. As explained in the section above on the Relative Sunspot Number, this is a worldwide center where sunspot data are collated and analyzed and the of fi cial International Sunspot Number worked out. Experienced observers can establish themselves as observing “stations” and contribute their results directly via the SIDC website – see http://sidc.oma.be for details.

Whether you just count active areas or measure R as well, it is a fascinating exercise to plot a graph of solar activity against time. To begin with, you can do this for a single year, plotting the number of active areas or R against the months on a simple line graph. You can plot graphs using graph paper, if you wish, but nowa-days it is also possible to draw them on a computer, using software such as the chart-making features on Microsoft Word or Excel. Unless the solar cycle is close to maximum or minimum, the resulting graph will almost certainly show an upward or downward trend, depending on which part of the cycle the Sun is in.

Monitoring sunspot activity levels over several years produces an even more interesting graph. Figures 5.5 and 5.6 show graphs of the MDF and R, respectively, plotted from the author’s own sunspot counts, for the decade beginning in January 2001 and ending in December 2010. Graphs plotted over several years have a jagged appearance, as activity can vary greatly from month to month, but both plots clearly show the latter half of solar cycle 23 and the beginning of cycle 24. Activity declines from near-maximum levels in 2001, goes through an extended minimum around 2008 and is quickly increasing again by the end of 2010. The shape of both graphs agrees quite well with the graphs produced by amateur solar groups from observa-tions by many workers – compare the plot of R in Fig. 5.6 with Fig. 5.7 , which shows

95The Relative Sunspot Number

0

2

4

6

8

10

12Ja

n-01

Jul-0

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MDF

Fig. 5.5 Graph plotted from the author’s sunspot counts, showing the variation in the MDF between 2001 and 2010

Fig. 5.6 Graph plotted from the author’s sunspot counts, showing the variation in the Relative Sunspot Number R between 2001 and 2010

96 5 Measuring Solar Activity

the collective graph produced by the BAA Solar Section for the same period. Many peaks and troughs in the graph agree as well – for example, in the author’s R graph the great peak of September 2001 shows up in the collective graph as well.

However, your personal graph can easily be distorted if you do not make many observations in a month. This can often happen in winter, when the few days on which you manage to make observations coincide with a period of unusually high or low activity. This gives an erroneous MDF or R fi gure for the month and causes a peak or trough on your personal graph that may be absent on the collective plot. An example of this happened to this author in March 2001. It was a very cloudy month, and there were only 3 days in which I was able to count the Relative Sunspot Number. Two of these clear days were at the end of the month, the time at which the very large sunspot group pictured in Fig. 3.9 (Chap. 3 ) was on the disc. This group contained many spots, giving very high R counts on both days. Activity earlier in the month, however, was much lower, and so my average R for the month was higher than it should have been. The peak visible in Fig. 5.6 does not show up in the collective BAA graph.

Graphs can be used to tell at a glance when the maximum or minimum of a solar cycle has passed – but only after the event. It is obvious from both graphs that sun-spot minimum occurred in mid-2008. But we must always beware of declaring too quickly when a maximum or minimum has passed, because the Sun can, and does, frequently surprise us. For example, at the time of the 2008 sunspot minimum, it was easy to assume that activity would soon start to pick up. In the event, long periods of little or no activity persisted for most of 2009, and it was not until 2010 that the new sunspot cycle took off. Only by carefully recording and plotting activity over several years can we be sure that a maximum or minimum has really passed.

0

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200Ja

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R

Fig. 5.7 Graph showing the variation in the Relative Sunspot Number R between 2001 and 2010, based on British Astronomical Association collective sunspot data (Permission to use BAA data kindly supplied by Lyn Smith, BAA Solar Section Director)

97Observing Faculae and White-Light Flares

Observing Faculae and White-Light Flares

As well as sunspots, you can also keep records of faculae. However, as explained in Chap. 3 , faculae are only easy to see when they are close to the limb. The best way of recording them is by drawing their positions and outlines on whole-disc drawings, in the same way as for sunspots, or by imaging them with a digital cam-era or webcam and, if necessary, increasing the contrast at the image processing stage. (Care must be taken when arti fi cially enhancing images, as this may affect their scienti fi c validity – see Chap. 8 .).

The most interesting type of faculae are the polar faculae, partly because of the challenge of seeing them and also because relatively little is known about them (see Chap. 3 for more details). To see these faculae you need a high-contrast solar image viewed at high magni fi cation in good seeing. Contrast is especially impor-tant, as polar faculae are low-contrast features, and if any daylight is allowed to fall onto the image your chances of seeing them are much reduced. You can look for polar faculae using a 9 mm (101×) eyepiece on an 80 mm refractor to project a 253 mm solar disc in your projection box, which gives good contrast, especially when you partly cover the open section. In fact, some of the best views of polar faculae occur when observing the Sun’s image in the projection box inside a curtained room, the curtains providing additional blackout. If you observe the Sun directly with a fi lter, use a rubber eye guard on the eyepiece to prevent stray light from reaching your eye. Some observers even shroud their entire face with a black cloth, as old-time photographers used to do when changing the plates. Contrast and seeing are more important than telescope aperture when observing polar faculae.

To fi nd polar faculae you need to know the values of P and B 0 for the day you

are observing. Faculae are always easiest to see at the pole that is tipped towards Earth. Therefore the best time to see faculae at the north pole is in autumn, as this pole reaches its maximum tilt towards us ( B

0 = +7.2°) in September. Similarly,

southern polar faculae are best seen in spring, B 0 reaching −7.2° in March. In

midsummer and midwinter, when B 0 passes through zero, it is sometimes possible

to see faculae at both poles. How many polar faculae you will see at either pole depends on the current stage of the solar cycle, as research has shown that num-bers of polar faculae vary inversely to the sunspot cycle, i.e., their numbers are highest at sunspot minimum and the declining phase of the cycle just before minimum.

To look for polar faculae, start at low power and ensure that the whole-disc image is correctly oriented using the spot-drift method. Then read off the value of P for the pole currently tipped towards Earth on the angular scale at the edge of the projection grid. This is the part of the limb you should next examine with the high-power eyepiece. Remember that polar faculae are not large patches or streaks, like spot-related faculae, but rather are very small, bright spots, sometimes little more than points of light (though some may appear slightly elongated because they are close to the limb).

98 5 Measuring Solar Activity

In addition to the contrast-enhancing measures mentioned above, you may fi nd the faculae easier to see if you move the telescope very slightly back and forth using the slow motion controls. This is because the human eye sometimes sees low- contrast objects better when they are moving. (Indeed, the same trick is used by deep-sky observers when searching for faint objects.) Also, if you use the projec-tion method, moving the image is a good way of assessing whether polar faculae are real and not just irregularities on the surface of the projection screen. Regardless of the value of B

0 , carefully scrutinize both poles for faculae, since faculae are

harder to see and much less frequent at the pole tipped away from Earth; even so, one or two are still sometimes visible. If you see any very bright or unusually shaped faculae at the poles, make a note of them and do a basic sketch or obtain some high-resolution images if you can.

The easiest way of making useful records of polar faculae is to simply note whether any are present each day you observe the Sun. Record at which pole or poles they are visible and also make a note of the seeing conditions at the time you are observing. At the end of the month you can calculate the percentage of observ-ing days on which you saw polar faculae and plot this fi gure on a graph against time. In fact, you could plot this on the same graph as, say, R , with a percentage scale on the opposite side to the Y-axis, to show how the frequency of polar faculae varies with the sunspot number. To make the statistics more accurate, you could include in your calculations only the days in which the seeing was at least reason-ably good. Counting the numbers of individual polar faculae is very dif fi cult, because the faculae are so small and elusive that they tend to drift in and out of visibility due to atmospheric turbulence, and their positions are dif fi cult to memo-rize. The short lifetimes of polar faculae add to the dif fi culty in counting them.

If you are ever fortunate enough to see a white-light fl are, you should report it immediately to your observing group coordinator by the quickest method possible, preferably by telephone. However, if you are relatively new to solar observing, contact an experienced observer at your local astronomical society and, if possible, get him or her to verify your observation. As noted in Chap. 3 , it is easy to mistake a light bridge or other types of bright areas within a sunspot for a fl are, and so there is the danger of causing a false alarm. If you see what you believe to be a white-light fl are, record the time to the nearest minute at which you fi rst noticed it, the time at which it reached maximum brightness and the time at which it faded from view. You should also, of course, record the details that you would for any solar observation, such as seeing conditions, instrument, method of observation and so on. Make a basic sketch of the fl are in relation to its parent sunspot group and, if possible, photograph it using the methods presented in Chaps. 7 and 8 .

Observing Naked-Eye Sunspots

A large sunspot group is sometimes visible to the naked eye, provided your eye is protected by a safe solar fi lter. Filters used for telescopic solar observing, such as Mylar or glass aperture fi lters, are just as effective for observing the Sun with the

99Observing Naked-Eye Sunspots

naked eye and, as mentioned in Chap. 2 , a welder’s glass is also suitable, provided it is of shade number 14. Whichever type of fi lter you choose, try to obtain one that is large enough for the Sun to be seen with both eyes at the same time. This is more comfortable than squinting with just one eye, and you will be able to see any naked-eye spots more clearly if your eyes are relaxed.

Observers have sometimes seen naked-eye sunspots without a fi lter, when the Sun is setting or is dimmed by mist or thin cloud. Indeed, in Chap. 1 we noted that the earliest recorded observations of sunspots were made in this manner by the ancient Chinese. However, while it is worth noting if you see a spot in this way by accident, it is not recommended as a method of deliberately looking for naked-eye spots. You can never be sure that the amount of solar radiation getting through to your eyes is small enough not to be dangerous.

A sunspot – or, more usually, a sunspot group – has to be at least fairly large to be visible to the naked eye. A rough rule of thumb used by astronomers for many years is that a group has to be at least three times the size of Earth in its longest dimension for it to be visible. But we have to take several important factors into consideration. For example, a spot must be well away from the limb for it to be seen with the naked eye, as when close to the limb its apparent size is greatly reduced by the effect of perspective. Also, the compactness of the group needs to be taken into consideration. A group containing large areas of umbra in a compact arrangement is much easier to see than a group of the same size whose spots are sparsely scattered. Finally, there is the observer’s eyesight. An observer with per-fect or near-perfect long-range vision can see smaller sunspots than someone with mediocre eyesight.

It is interesting to check each day whether any naked-eye sunspots are visible, and some solar observing organizations keep systematic records of naked-eye sunspots as a way of monitoring solar activity. To do this, examine the Sun with your fi lter once each day and note down the number of sunspots you can see. Always look for naked-eye sunspots before observing the Sun with a telescope. If you see a large sunspot with your telescope fi rst, your chances of seeing it unaided will be much higher, because you will know it is there and will know its position on the Sun’s disc as well. This observer bias is sometimes unavoidable, as you may well know that a large spot is present because you saw it using the telescope the previous day.

Usually the number of naked-eye sunspots seen each day will be 0 or 1. At solar minimum, looking for naked-eye sunspots can be patience-testing, as there will be periods of weeks or months when no naked-eye groups are visible. Near maximum, however, naked-eye spots are quite common, and there are sometimes days when two or more spots are visible on the disc at the same time. At the end of each month, add up the number of spots and average them to get an MDF, as for sunspots seen using the telescope. You can then plot the MDFs on a graph to show how naked-eye sunspot numbers vary over time. Past observations indicate that naked-eye sunspot numbers correspond quite closely with the MDF and R as seen through telescopes, so that the maximum number of naked-eye sunspots occurs roughly at the time of telescopic sunspot maximum.

101L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0_6, © Springer Science+Business Media New York 2012

Observing the Chromosphere

Chapter 6

As noted in Chap. 1 , the Sun is surrounded by an extensive, invisible atmosphere. This is in two parts: the chromosphere, a thin layer a few thousand kilometers thick, and the corona, which consists of rare fi ed but extremely hot gas extending millions of kilometers out into space. Both layers are invisible to the Earth-bound observer in ordinary circumstances, because the brilliant light of the photosphere is scattered by Earth’s atmosphere and so washes out the relatively feeble light of the solar atmosphere.

The only time an Earth-based observer can see the Sun’s atmosphere without special equipment is during the fl eeting few minutes of a total solar eclipse. To see the corona from Earth without waiting for an eclipse is very dif fi cult, due to its extreme faintness relative to the Sun’s surface. Professional astronomers can pho-tograph the corona using an instrument called a coronagraph , a telescope contain-ing an occulting disc to precisely block out the brilliant photosphere and optics specially made to keep scattered light from the photosphere to an absolute mini-mum. For a coronagraph to work, it must be installed at a very high mountaintop site, where the scattering of light by our own atmosphere is minimized. Even then, a coronagraph can only image the inner parts of the corona. The best images of the corona are obtained from space-based solar observatories.

The chromosphere, however, is a different matter. As discussed at the beginning of Chap. 2 , visible sunlight is made up of a spectrum of colors, running from violet light at a wavelength of 400 nm to red at 700 nm. The photosphere and corona emit continuous spectra – that is, their light is emitted evenly across the whole range of visible wavelengths – but the chromosphere emits most of its light in a number of fi ne lines, each of them covering a very narrow band of the visible spectrum. If we can isolate one of these lines and observe the Sun in just this

102 6 Observing the Chromosphere

wavelength, most of the brilliant light of the Sun – and consequently the scattering by Earth’s atmosphere – would be eliminated, and we would see just the light of the chromosphere.

The brightest line emitted by the chromosphere is known as hydrogen-alpha (abbreviated H-alpha or H a ) and is located in the red part of the spectrum, at a wavelength of 656.3 nm (6563 Å) 1 . Another strong chromospheric emission line is calcium-K (abbreviated Ca-K), at the extreme violet end of the spectrum, wave-length 393.4 nm (3934 Å). The chromosphere also emits in the blue light of hydro-gen beta as well as helium (yellow) (Fig. 6.1 ). These wavelengths – and many more – are studied by professional astronomers, but H-alpha and Ca-K are much the easi-est to observe, and most instruments made for the amateur are designed to show the Sun in these wavelengths.

Astronomers use two basic types of instruments to isolate the H-alpha or Ca-K spectral line. The oldest method involves splitting up the solar spectrum to a suf fi cient extent that an individual emission line is resolved and isolated. A more recent development is a special type of fi lter known as an interference fi lter, which cancels out all wavelengths of sunlight reaching the observer except, say, H-alpha. Because these fi lters show the Sun in only one color at a time, interference fi lters for observing the chromosphere are known as monochromatic fi lters. These fi lters are nowadays available commercially, at prices amateurs can afford. It is wonderful that the amateur can see the chromosphere, for in this layer we can see many features that are not observable in white light and begin to see some of the activity revealed so dramatically in images taken by professional and space-based observatories.

When we observe in H-alpha, the Sun’s disc has a brilliant red color and the chromosphere appears as a thin, faint band around the limb. The “surface” – i.e., the chromosphere – has a granular structure, but this is much larger and easier to see than the photosphere granulation seen in white light. Sunspots are still visible, but they are not as prominent as they are in white light, because we are looking at them through a welter of activity in the chromosphere.

4000A° 5000A° 6000A° 7000A°

Ca-K3934A°

Violet

H-α6562.8A°

Red

Ultraviolet Infrared

Fig. 6.1 Diagram of the visible light spectrum, showing the wavelengths of two major chro-mospheric emission lines (Diagram by Dominic Ford)

1 1 Å unit = 0.1 nm or 10 −10 m.

103 Observing the Chromosphere

Four major solar activity features are visible, all of them associated with mag-netic fi elds and broadly following the solar cycle in their numbers and intensity. The most dramatic H-alpha features are the prominences – large, fl ame-like structures of hydrogen at the limb of the Sun. Many of these are tens of thousands of kilometers high and project out into what looks like the blackness of space but which is, in fact, the invisible corona.

Prominences vary hugely in size, shape and structure. Some appear as small spikes at the limb of the Sun; others are intricate loops appearing exactly like the arrangement of iron fi lings on paper above a bar magnet – revealing that they are caused by magnetic fi elds. Another type extends to only a low height in the solar atmosphere but covers a very large area of the limb. Some types appear to erupt out of the Sun like lava from a volcano and change structure very rapidly. Changes in prominences do not normally occur fast enough for them to take place before your very eye, because at the Sun’s distance material being ejected at very high speeds appears to be moving quite slowly.

Prominences are often associated with sunspots, but not always, and sometimes when sunspot activity is quite low it is surprising how many prominences are visi-ble around the limb in H-alpha. During such periods of low sunspot activity there is sometimes a large and spectacular prominence lurking in this wavelength. This picture of a huge prominence (in Fig. 6.2 ) was taken on a day when white-light

Fig. 6.2 Example of how the level of white-light sunspot activity is no accurate guide to what might be visible in H-alpha. This giant solar prominence was photographed by the author on April 4, 2004, when sunspot activity was very moderate (see text). Picture taken using a Canon EOS 300D digital SLR camera attached to an 80 mm refractor with Baader “H-alpha corona-graph” prominence viewer. Exposure 1/250 s with the camera set to ISO 400

104 6 Observing the Chromosphere

activity was very moderate: there were just four active areas on the disc and the Relative Sunspot Number R measured only 64. With limited time for observing that morning and the unremarkable level of activity, it seemed to warrant only to have a cursory look in H-alpha. This experience is a dramatic illustration of the fact that the level of white-light activity is not an accurate guide to what might be visible in special wavelengths.

If you observe the Sun with a suitable H-alpha instrument on almost any day you will probably notice at least one or two sinuous, dark lines on the disc that are not visible in white light (Fig. 6.3 ). These are called fi laments and are simply prominences seen silhouetted against the disc. More strictly, we should say that they are seen in absorption , as they are cooler and less dense than the chromo-sphere beneath them and so absorb some of the light from the chromosphere. It is sometimes possible to see their true nature yourself when a fi lament close to the limb extends into the darkness above the limb, and part of it shows up as a promi-nence. Such a combined fi lament and prominence is sometimes known informally as a “ fi laprom” (Fig. 6.4 ). Like prominences, fi laments are often associated with sunspots and sometimes appear to be entangled with sunspot groups, but on many occasions they occur well away from any spots and at high latitudes, where spots seldom or never form.

Fig. 6.3 The Sun’s disc in hydrogen-alpha. The dark linear features are fi laments (i.e., promi-nences seen in absorption against the disc), and the bright regions are plages. Whole-disc mosaic image by Dave Tyler

105 Observing the Chromosphere

If there is any activity at all on the Sun, you will quite likely notice some large, bright patches, many of them surrounding sunspots. These are known as plages or fl occuli and resemble the faculae seen around sunspots in white light, except that unlike the latter they are visible at the center of the disc and not just when near the limb. Like the faculae, they often form several days before the sunspots themselves and linger for some time after the spots have disappeared. Many plages cover a considerable area and are a graphic illustration that a sunspot group is often just the “tip of the iceberg” in a large active region.

Finally, with an H-alpha fi lter solar fl ares become quite common features, several a day sometimes being visible at solar maximum. A flare shows up as a brilliant white patch which suddenly appears in the space of a few minutes. Although fl ares reach their peak brightness very quickly, they can take an hour or more to fade from view and can sometimes resurge again, extending the spectacle for much longer.

If you observe in Ca-K, the view is slightly different from H-alpha. Plages are still visible – indeed, they are the most prominent feature at this wavelength – and tracking plages is the most useful type of observation to be made at this wavelength. Filaments, however, are generally not visible, although it is possible to see (and image) some of the brighter prominences. The chromospheric granulation is much in evidence, though this shows a different structure from that seen in H-alpha. But observing in Ca-K poses a major dif fi culty: you may not be able to see the Sun at all!

Fig. 6.4 An example of a “ fi laprom” – a fi lament seen in absorption against the solar disc extending beyond the Sun’s limb and becoming a prominence seen in emission against the apparent blackness of space (In reality, it is seen against the corona, but the latter is much too faint to be seen in an H-alpha fi lter.) (Image taken by Dave Tyler on November 20, 2009)

106 6 Observing the Chromosphere

As we can see from the diagram in Fig. 6.1 , Ca-K is located in the violet, at the very edge of the visible spectrum. How far you can see into the violet varies from person to person, and some people with good sensitivity at these wavelengths will see the calcium Sun with no dif fi culty and enjoy a fi ne view. But others may fi nd that they can see the Sun only very faintly, or even not at all, because the limit of their optical sensitivity in the violet falls short of the 393.4 nm wavelength of Ca-K. In particular, human sensitivity to these short wavelengths decreases with age, and many people over 50 years old may have problems seeing the Sun in Ca-K. Sometimes it helps to wait a minute or two for your eyes to adapt to the darkness, much as one would wait for the eyes to dark-adapt before enjoying a good view of the night sky. It also may be helpful to use some sort of Sun-shield around your eyes when observing such a faint solar image – indeed, Sun-shields are sometimes avail-able commercially with solar telescopes and fi lters.

Observing in Ca-K poses a further problem in that when the Sun is at a low altitude, violet light and other short wavelengths are scattered by Earth’s atmo-sphere, and only the longer wavelengths, i.e., orange, red and yellow, pass through to the observer’s eye. (This is why the Sun appears red when it is rising or setting.) So observing in Ca-K is best done when the Sun is at a high altitude, preferably 40° or higher above the horizon.

If you live at a high temperate latitude such as the northern United States, Canada, or northern Europe you may fi nd that observing in calcium light is restricted to the summer months and even then may only be possible in the middle of the day. (This altitude problem works in reverse in H-alpha: because H-alpha is located deep in the red part of the spectrum, this wavelength is hardly absorbed at all by the atmosphere when the Sun is at low altitude, and you can enjoy spectacular views of H-alpha features almost down to sunset.) However, imaging in Ca-K is much less problematic than visual observing, as many digital cameras and web-cams have excellent sensitivity at violet wavelengths. The only problem is in seeing the image on the view fi nder in order to center and focus it, so Ca-K imaging is best done on a computer screen rather than by direct viewing at the camera.

When buying a Ca-K fi lter system or dedicated Ca-K telescope (see below), think carefully about what you want to use it for. If you are primarily interested in visual observing at this wavelength, it is a good idea to have a look through some-one else’s Ca-K fi lter (perhaps at an astronomy club meeting) before parting with your money, to make sure you can see the Sun’s image easily. Ca-K systems are very specialized instruments, best suited to those interested in digital imaging of plages, which are best seen at this wavelength. H-alpha is much the easiest and most versatile wavelength for the amateur solar observer, and the bulk of this chap-ter will concentrate on H-alpha fi lter systems and what can be done with them.

Equipment for Observing the Chromosphere

Unfortunately, the equipment required to isolate the H-alpha (and Ca-K) lines and reveal the chromosphere is expensive. A typical H-alpha fi lter costs as much as a very good telescope, and this is the reason why, until recently, only a minority of

107Equipment for Observing the Chromosphere

amateur astronomers observed in H-alpha. However, in recent years technological advances have brought the price of these fi lters down enormously, so that fi lters at the entry-level end of the market are now within reach of the average amateur. Also, just as in white light, you can do H-alpha work with a small telescope.

So that we can understand how H-alpha fi lters work and assess what is the most suitable equipment to buy, it may be helpful to brie fl y look at the history of H-alpha solar observing. Before the mid-nineteenth century astronomers had to wait for a total solar eclipse to see the Sun’s atmosphere and prominences. Indeed, some astronomers thought that prominences were part of the Moon and not the Sun at all, and it took careful use of the new medium of photography at the eclipse of July 1860 to con fi rm that prominences are solar in origin.

The fi rst observation of prominences outside a solar eclipse was made indepen-dently in 1868 by the British astronomer Norman Lockyer (1836–1920) and the Frenchman Pierre Janssen (1824–1907), by aligning the slit of the recently invented spectroscope at a tangent to the limb of the Sun and thus isolating the H-alpha spectral line. By opening up the spectroscope slit a little it was possible to see a prominence, if one was present on that part of the limb. Observing prominences in this way was rather cumbersome, though, because it meant viewing prominences one at a time by patiently scanning around the limb. But until the advent of afford-able H-alpha fi lters almost a century later, the prominence spectroscope remained the standard method of viewing the prominences for amateur astronomers – and those amateurs that had spectroscopes were a privileged few.

In the 1890s an improvement on the spectroscope, known as the spectrohelio-graph , was invented by the American George Ellery Hale (1868–1938) and another Frenchman, Henri Deslandres (1853–1948). This allowed astronomers to photograph larger areas of the Sun, including disc features in the chromo-sphere such as fi laments and fl ares, as well as prominences. In the 1920s Hale made a modi fi ed version of this instrument, known as the spectrohelioscope , which allowed the disc features to be viewed visually. The spectroheliograph remains a standard instrument in professional observatories to this day, and a small number of amateur astronomers observe the Sun using home-made spectrohelioscopes.

The spectrohelioscope has a major advantage over H-alpha fi lters in that it can be “tuned” to any wavelength in the spectrum and thus show the chromosphere in any of the emission lines, including calcium and hydrogen-beta as well as H-alpha. An observer with a spectrohelioscope can thus compare the view of chromospheric phenomena in different wavelengths. However, spectrohelioscopes are not available commercially and require a great deal of optical and mechanical skill to build, put-ting them outside the reach of most amateurs. In addition, unlike fi lters, they are very bulky and need to be set up in permanent observatories.

The development of fi lters for H-alpha observing began in 1930, when yet another Frenchman, Bernard Lyot (1897–1952), invented the coronagraph and made the fi rst non-eclipse observations of the corona from the high-altitude site of Pic du Midi Observatory in the French Pyrenees. In 1939 Lyot improved the coro-nagraph by incorporating an interference fi lter – a complex fi lter composed of many layers of quartz and polarizing materials that, using the principle of interference,

108 6 Observing the Chromosphere

canceled out all wavelengths of solar radiation except the H-alpha line. This – the world’s fi rst H-alpha solar fi lter – enabled spectacular images of prominences to be taken through the coronagraph. Lyot’s fi lter could be used to observe other spectral lines as well. Interference fi lters of Lyot’s design were later improved to the point where they could show disc features as well.

To show chromospheric features at all, a fi lter (or spectrohelioscope) must isolate a tiny slice of the Sun’s spectrum. Light at the red end of the visible spec-trum has a wavelength of around 700 nm (7,000 Å), while at the violet end the wavelength is about 400 nm (4,000 Å). The visible spectrum, therefore, covers about 300 nm (3,000 Å), but the H-alpha line is only about 0.1 nm (1 Å) wide. The extent of the spectrum allowed through by a fi lter is known as the fi lter’s passband. The narrower the passband, the more the H-alpha line is isolated from the main part of the Sun’s light and so the greater the contrast of the resulting images. Filters with narrower passbands are more complex, and therefore more costly, to make.

A fi lter with a passband of several angstroms will reveal prominences, provided it is centered on the H-alpha line, the Sun is hidden by an occulting disc and the sky is free from contrast-reducing haze. But to reveal features on the disc it is necessary to use a “sub-angstrom” fi lter – i.e., one with a passband of less than 1 Å, narrower than the H-alpha line itself. Filters of the type employed by Lyot, using quartz and Polaroids to isolate the H-alpha line, are still used by professional astronomers today as a standard means of imaging the chromosphere, but they are too bulky to fi t on amateur telescopes, and their fi ve- or six- fi gure price tags put them well beyond the amateur’s budget.

In the 1960s a new type of interference fi lter became available. Instead of quartz and Polaroids it used a large number of layers of dielectric materials such as mag-nesium fl uoride deposited on glass, known as an etalon , to isolate the H-alpha line. These fi lters were much more compact and, most importantly, were available “off the shelf” at prices amateurs could afford. Initially they were only capable of show-ing prominences, because their passbands were wider than 1 Å, but a number of amateurs successfully developed prominence-viewing devices using the new fi lters mounted in telescopes built on the coronagraph principle, with a disc to occult the Sun. The following decade the U.S.-based DayStar Filter Corporation developed sub-angstrom fi lters of the same type, enabling amateurs for the fi rst time to observe disc features as well as prominences with a fi lter attached to an ordinary telescope.

Because the market for such specialized fi lters is very small, only a handful of companies around the world currently produce them. Nevertheless, the amateur astronomer today can choose from a large and growing number of commercially available H-alpha fi lter systems. Indeed, the range of systems available can seem quite confusing and, moreover, is changing rapidly, so that by the time this volume appears in print some of the telescopes and fi lters described below may have been discontinued and new models may have appeared.

To observe the Sun in H-alpha you can either choose a complete H-alpha solar telescope (which cannot be used for anything else) or obtain an H-alpha fi lter with

109Equipment for Observing the Chromosphere

a mount to adapt it to an existing telescope, normally a small refractor. Most H-alpha fi lters on the amateur market today are of the sub-angstrom variety, although some entry-level models have passbands barely below 1 Å. All mono-chromatic fi lters are expensive by ordinary standards, but they are a good invest-ment, considering the wealth of otherwise invisible solar detail they reveal.

For an interference fi lter to work, the light rays striking it must be close to parallel. Monochromatic fi lters come in two types: those in which the main interference fi lter – the etalon – is mounted in front of the telescope objective, and those in which the etalon is at the eyepiece end. Basically, in fi lter systems made by Coronado, Lunt Solar Systems and Solarscope the etalon is at the front of the tele-scope, whereas DayStar fi lters are mounted at the rear. In fi lters mounted in front of the objective, the unfocused solar rays are already parallel, so this is not a problem. But as fi lters made by DayStar (and one or two other, less well-known manufacturers) are mounted at the eyepiece end of the telescope, the light rays striking them are convergent. In most amateur telescopes, whose f/ratios are rarely longer than f/15 and often much shorter, this convergence is very strong, and it has the effect of widening the passband of the fi lter and so reducing the amount of detail visible through it. H-alpha fi lters work best if the effective f/ratio of the telescope is increased to about f/30, so that the light rays are closer to parallel. With DayStar-type fi lters, you need to use a pre- fi lter or energy rejection fi lter (ERF) mounted in front of the telescope’s aperture to achieve this. The ERF increases the f/ratio by stopping down the aperture, because it is normally much smaller than the tele-scope’s aperture. It also contains a fi lter made of heat-resistant red glass to protect the main H-alpha fi lter (and your eyes) from the solar heat.

Never use an H-alpha fi lter without an ERF: you could risk damaging your eye-sight, as well as the fi lter. Because the size of the ERF varies with the aperture and f/ratio of your telescope, you need to specify these details when ordering a DayStar-type fi lter. However, with many of today’s short-focus refractors, stopping down the telescope may not be enough to achieve an f/30 light cone – or, looking at it another way, stopping down to f/30 would produce too small an aperture for useful detail to be resolved on the Sun’s image. In this case, as well as the ERF, you need to use a special “telecentric” lens in front of the etalon; this increases the effective f/ratio of your telescope. To fi nd out the correct lens you need, you should contact DayStar or a DayStar dealer at the time of ordering your fi lter.

It is not a good idea to insert an ordinary Barlow lens in front of the etalon, because a Barlow causes the light rays to diverge and so will compromise the per-formance of the etalon, which requires the incoming light rays to be as close as possible to parallel in order to work to its full performance. (One exception to this rule is the “Powermate” range lenses made by Tele Vue; these increase the f/ratio without causing the rays to diverge.) In front-etalon fi lters such as those made by Coronado, the etalon is mounted with the ERF and forms part of the unit that goes over the telescope objective. In addition to the H-alpha line, all etalons transmit “side bands” of small parts of the spectrum at other wavelengths, and these have to be eliminated by a “blocking” fi lter behind the etalon. In DayStar fi lters the blocking fi lter is located directly behind the etalon, but in the front-etalon fi lters it is located

110 6 Observing the Chromosphere

at the eyepiece end of the telescope, often forming part of a star diagonal that re fl ects the solar image through 90°.

H-Alpha Telescopes

The cheapest H-alpha fi lter systems are, in fact, complete H-alpha solar tele-scopes. The best known of these is the 40 mm aperture “Personal Solar Telescope,” or PST, made by Coronado (part of Meade Instruments Corp.), introduced in 2004 (Fig. 6.5 ). The PST has an internal H-alpha fi lter system with a passband of slightly less than 1 Å – rather wider than other Coronado fi lters – but it is still a sub-angstrom system, narrow enough to reveal a signi fi cant amount of disc detail as well as prominences.

Because the effectiveness of an interference fi lter depends very sensitively on the angle the light rays are hitting it (see above), all H-alpha fi lters, including those built into the PST, need to be “tuned” very precisely to the H-alpha line by tilting them back and forth using a built-in thumbscrew – in the case of the PST, by turning a collar half way along the tube.

Using modern digital cameras and webcams, many amateurs have taken spectacular images with these instruments, showing remarkable detail for a system

Fig. 6.5 The Coronado Personal Solar Telescope (PST) – the 40 mm dedicated H-alpha tele-scope that made sub-angstrom H-alpha views of the Sun available for under $500. The large dark sheet is a Sun shade to shield the observer from direct sunlight (Photograph by the author)

111Equipment for Observing the Chromosphere

that has a theoretical resolution of about 4 arc sec. Perhaps the most delightful feature of the PST, and other small H-alpha telescopes, is its extreme portability. It is small and light enough to be mounted on an ordinary camera tripod and taken out at a moment’s notice, making it a very attractive option if you live in a cloudy climate. It can also be easily mounted “piggyback” on an existing telescope, perhaps an instrument showing the Sun in white light at the same time, thus enabling you to directly compare the white-light and H-alpha views.

Inevitably there are compromises in an H-alpha telescope at its very reasonable price. The contrast of the image is not as high as that in more expensive models, the mechanical controls (focusing and tuning) are very basic and some observers have complained of distracting ghost images and a “sweet spot” in the fi eld of view, in which H-alpha features show up with better contrast than elsewhere in the fi eld. You can greatly increase the contrast of the disc features in a PST (and other Coronado fi lter systems) by purchasing an extra etalon from Coronado and mount-ing it in front of the existing etalon, an arrangement known as “double-stacking.” However, this automatically doubles the price of the system, and it also makes the prominences fainter. If you have the money to buy an extra etalon, you might want to consider buying one of the more expensive models of H-alpha telescope or fi lter. Nevertheless, the Coronado PST offers remarkable value for money and is an excel-lent choice if you are new to H-alpha observing or your budget is restricted.

More recently, the German fi rm Lunt Solar Systems has also introduced an entry-level H-alpha telescope at an attractive price. The Lunt LS35THa in its most basic form is of a similar design to the Coronado PST: etalon at the front, blocking fi lter in a star diagonal at the rear. Its aperture is even smaller than the PST: just 35 mm, but some good results have been claimed for it, and it may well become a European competitor to the PST. Like the PST, it is extremely compact and porta-ble, making it usable at a moment’s notice.

Both Coronado and Lunt make larger, more expensive H-alpha telescopes. For several years, Coronado has offered its “SolarMax” range of 40, 60 and 90 mm apertures, whose fi lters are guaranteed to have passbands of less than 0.7 Å. More recently, Coronado has introduced the “SolarMax II” 60 mm and 90 mm H-alpha telescopes, which in some cases cost less than the equivalent telescopes in the older SolarMax range. The main difference between the two sets of models is the way in which the fi lters are tuned to the H-alpha line. The older models use the “T-Max” tuner, which is essentially a hinged arrangement in which the etalon is mounted on the front half of the “hinge,” and the alignment of the fi lter can be adjusted very delicately by a thumbscrew. The SolarMax II telescopes use a tuning arrangement called “RichView,” which allows the etalon to be tuned to either side of the H-alpha line. The T-Max tuner allows the fi lter to be shifted slightly away from H-alpha to a very slightly shorter wavelength. As explained below under “Prominences and Filaments,” this allows you to see erupting fi laments moving toward Earth. The RichView tuner allows you to see material falling back down toward the Sun as well. All these telescopes, including the PST, are dedicated H-alpha solar tele-scopes and cannot be used for any other purpose.

112 6 Observing the Chromosphere

As in the PST, with both SolarMax and Solar Max II telescopes it is possible to narrow the passband, and thereby increase the contrast of solar disc features, by “double-stacking” an additional etalon in front of the existing one. Again, this does signi fi cantly increase the cost. For a smaller increase in price, you can order a larger blocking fi lter than the standard one that comes with the telescope. This allows a bigger fi eld of view as you see it through the eyepiece – a larger fi eld allows you to see the whole Sun in the fi eld at a decent magni fi cation and allows cameras with larger imaging sensors to be used.

Lunt Solar Systems sells several H-alpha telescopes (Fig. 6.6 ), which roughly parallel the Coronado systems in aperture, price and performance: 60, 80 and 100 mm refractors are available in a wide range of prices. It also offers two “giant” H-alpha telescopes, a 152 mm and the most recent addition a whopping 230 mm dedicated H-alpha telescope, at a suitably whopping price! As with Coronado, all the Lunt telescopes have passbands guaranteed to be narrower than 0.7 Å.

Different models of Lunt telescopes are available in the same aperture at a wide range of sophistication, and for a given aperture the price you pay for a Lunt tele-scope varies greatly according to the speci fi c model you buy. At the bottom of this scale is the basic 60 mm H-alpha telescope (model LS60THa), with a standard-sized blocking fi lter in the star diagonal. The more expensive models offer a larger blocking fi lter. Also, instead of the hinged mechanism on the older Coronado instruments and the Lunt 35 mm model, the higher-priced Lunt 60 mm refractors, and all their larger telescopes, use a “pressure tuner” to tune the fi lter to H-alpha.

Fig. 6.6 The Lunt Solar Systems 152 mm (6 in.) dedicated H-alpha solar telescope. In the background are Lunt’s 50 and 100 mm models (Photograph by the author)

113Equipment for Observing the Chromosphere

A “pressure tuner” works by slightly changing the air pressure in the chamber containing the etalon, which changes the refractive index of the air inside the cham-ber and thus varies the fi lter’s peak spectral sensitivity. Like Coronado’s “RichView” system, pressure tuning allows you to tune the fi lter to both the red and the blue sides of the H-alpha line, thus enabling you to see Doppler-shifted material moving towards or away from you.

Solarscope, based on the Isle of Man in the British Isles, offers two dedicated H-alpha telescopes in 50 mm (Fig. 6.7 ) and 60 mm apertures. Both telescopes have the compact and ultra-portable look and feel of the Coronado PST, but with larger apertures. They are also signi fi cantly more expensive even than telescopes of equivalent aperture from Coronado and Lunt. But Solarview prides itself on the superior quality of its telescopes and fi lters; in particular, it cites the lack of a central obstruction that is a feature of some competing fi lters (the central spot allows fi lters to be made more cheaply), and it claims better consistency of image quality across the fi eld of view, with no “sweet spot” like that present in some other H-alpha systems. The images are indeed better than other systems of similar aperture – though, of course, the price is much higher. As is the case so often, “you get what you pay for.” As with Coronado and Lunt, the Solarview telescopes have passbands of 0.7 Å, but this can be narrowed to 0.5 Å by double-stacking with an extra etalon.

Much the oldest source of sub-angstrom H-alpha fi lters for amateur astronomers is the U.S. fi rm DayStar, originally based in California, now located in Missouri. In addition to a full line of fi lters, DayStar now produces a dedicated 60 mm H-alpha

Fig. 6.7 The Solarview 50 H-alpha telescope made by Solarscope Ltd. (Photograph by the author)

114 6 Observing the Chromosphere

telescope, known as the SolaREDi. Its focal length, 1,375 mm, is somewhat longer than competing 60 mm H-alpha telescopes. It is available in three different pass-bands: 0.7 Å, 0.5 Å and an incredibly narrow 0.3 Å. In some models, the telescope is tuned to H-alpha by tilting the etalon; in others, tuning works by controlling the etalon’s temperature with a thermostat, as in some of the more expensive DayStar fi lters (see below).

H-Alpha Filters

All four of the above companies – Coronado, Lunt, Solarscope and DayStar – also produce H-alpha fi lters that you can mount onto your existing telescope. In the larger sizes, fi lters are a cheaper option than buying a dedicated H-alpha telescope – but, ironically, the 40 mm aperture SolarMax fi lter produced by Coronado (Fig. 6.8 ) is substantially more expensive than the PST, despite the latter having the same aperture. This is because the 40 mm SolarMax has a nar-rower passband (0.7 Å) and is generally aimed at the more serious solar observer than the entry-level PST. Nevertheless, the “SolarMax 40” was a breakthrough when fi rst introduced in 2001, because it was the fi rst sub-angstrom H-alpha fi lter to be available for under $1,000. This author has successfully used a SolarMax 40 since 2002 – fi rst with a Meade ETX telescope, then a Vixen 80 mm achromatic refractor and, since 2006, with a Takahashi FS-60C 60 mm

Fig. 6.8 A 40 mm Coronado SolarMax fi lter fi tted to the author’s Takahashi FS-60C 60 mm apochromatic refractor. (Photograph by the author)

115Equipment for Observing the Chromosphere

apochromatic refractor mounted piggyback on the Vixen 80 mm (which is now used mostly for white light solar observing). Its performance has been excellent on all three telescopes, though with the ETX the fi eld was too narrow for the whole Sun to be visible at once (see below). Prominences stand out beautifully against a very dark background. Filaments and the chromospheric granulation show up with plenty of contrast, as do plages around active regions and the occasional solar fl are.

Only the larger sunspots are visible underneath all this chromospheric detail, and it is dif fi cult or impossible to pick out the smaller ones, but the T-Max tuning device (the same one as supplied with the complete Coronado telescopes) allows the fi lter to be tilted far enough off the H-alpha line for the chromosphere and its features to disappear, and the Sun becomes just a red disc, with sunspots showing up almost as they do in white light. The fi lter does have a “sweet spot,” in which prominences and other H-alpha features jump into greater contrast, but it is easy to remember its position in the fi eld of view and get used to it.

For comfortable solar observing, you can use a star diagonal in a horizontal posi-tion with the eyepiece pointed to the left, and in this position the sweet spot is near the lower left edge of the fi eld of view. When counting prominences for the purpose of measuring solar activity, always carefully scan the entire solar limb through the sweet spot, using the telescope’s slow motions, because although the whole Sun and prominences are within the fi eld at the magni fi cation (35×), the contrast is somewhat lower in other parts of the fi eld and fainter prominences could easily be missed, leading to an erroneous prominence count. The 35× seems to be about the best magni fi cation with the SolarMax 40. Lower magni fi cations would give a brighter image, but at the cost of losing fi ne detail, while there is no advantage in using a signi fi cantly higher magni fi cation, because the image becomes faint and harder to see in daylight. In any case, the image produced by the fi lter’s 40 mm aperture has a resolution of only about 4 arc sec, so there is no advantage in using a high magni fi cation.

As with their complete solar telescopes, Coronado fi lters come in the older SolarMax format with T-Max tuners and in the new SolarMax II versions with RichView tuners. As noted below, it is important to choose the correct blocking fi lter for your telescope’s focal length. You can use a Coronado fi lter on a telescope whose aperture is larger than that of the fi lter; for example, a 40 mm SolarMax on a 60 mm refractor. However, to obtain the maximum resolution possible in your telescope with a Coronado fi lter, the etalon needs to have the same diameter as your telescope’s aperture. Therefore the larger your telescope the bigger – and more expensive – the fi lter needs to be if you are to make the most of your telescope’s aperture. Prices quickly rise with the size of the etalon. In addition to their 40 mm range, Coronado offers Solar Max II fi lters in 60 and 90 mm sizes. As with their complete H-alpha telescopes, the passband of Coronado fi lters can be narrowed from 0.7 to about 0.5 Å by double-stacking with an extra etalon.

Coronado offers three eyepieces and a basic (achromatic) 2× Barlow lens speci fi cally optimized for solar observing with their H-alpha fi lters. The deep red light of H-alpha has slightly different properties from white light, so optics that are

116 6 Observing the Chromosphere

optimized for H-alpha should in theory give somewhat better images. Marketed under the trade name “CEMAX,” these eyepieces have focal lengths of 12, 18 and 25 mm and look very attractive with their brass fi nish. They also come with rubber eye guards, which are pretty much essential when observing in sunlight. With a SolarMax 40 fi lter, there is barely any difference in performance between the CEMAX eyepiece and ordinary eyepieces for night-time astronomy, but a larger-aperture fi lter might offer a more sensitive test. In any case, the eyepiece is excel-lent for the job and has just the right “eye relief.” 2 When the eye guard is removed, the eye end of CEMAX eyepieces is identical to the popular “Series 4000” Plössl eyepieces manufactured by Meade, which enables you to attach a digital camera to them with the appropriate adapter (see Chap. 7 ). The 2× CEMAX Barlow gives too high a magni fi cation for it to be useful on the SolarMax 40, but it is very useful for night-time astronomy, as are the CEMAX eyepieces themselves.

At present, Lunt offers a series of H-alpha fi lters roughly equivalent to their complete H-alpha telescopes. The fi lters are available in four sizes: 35, 50, 60 and 100 mm. Their passbands are advertised as being less than 0.7 or 0.75 Å, which can be narrowed to less than 0.5 or 0.55 Å by double-stacking. Lunt fi lters take the same general pattern as Coronado: the etalon is mounted over the telescope objec-tive and the blocking fi lter is housed in a star diagonal in front of the eyepiece.

Solarscope also does a range of fi lters in four apertures: 50, 60, 70 and 100 mm. As with their telescopes, Solarscope fi lters are aimed at the high end of the market and are very costly for their apertures. These fi lters all have passbands of 0.7 Å; as with Coronado and Lunt, double-stacking reduces the passband to 0.5 Å at signi fi cant extra cost. Like the Solarview telescopes made by the same company, they have no central obstruction. Unlike Coronado and Lunt fi lters, the blocking fi lter in the Solarscope models is not in a star diagonal, but is simply a “straight through” fi lter assembly that you insert just in front of the eyepiece. This requires much less focus travel than a star diagonal and is an important consideration if photography is a high priority, as depending on your telescope, some cameras may not reach focus when mounted on the end of a star diagonal. Also, mounting the camera on a star diagonal can cause problems with steadiness, particularly if your camera is heavy, as then the camera is well outside the center of gravity of the optical system.

DayStar has been producing sub-angstrom H-alpha fi lters since the 1970s, and until the late 1990s they were the only fi rm producing sub-angstrom fi lters for the amateur market. As noted above, in DayStar fi lters the etalon is mounted at the rear of the telescope, just in front of the eyepiece, and only the ERF is placed in front of the telescope objective. The ERF can be purchased separately from the main part of the fi lter and is much the cheapest part of a DayStar fi lter setup. This gives DayStar fi lters a major advantage over other makes: if you upgrade to a larger

2 Eye relief is the distance behind the eyepiece where you need to put your eye to see the whole fi eld. If the eye relief is too small, you need to place your eye uncomfortably close to the eyepiece and may not be able to see the whole fi eld at all if you have to wear glasses. If the eye relief is too large, the fi eld becomes “vignetted” (cut off at the edge) if your eye is too close; moving your eye too far away can cause ambient daylight to interfere with your view of the Sun.

117Equipment for Observing the Chromosphere

telescope, you don’t have to spend a lot of money on a whole new fi lter system – you just need to get a larger ERF. For the same reason, it is possible to use the same DayStar fi lter on more than one telescope, as long as you have the appropriate ERF for the various telescopes. Also, the focal length of your telescope does not matter, because the blocking fi lter is the same size as the etalon and is more than large enough to show the full fi eld with most eyepieces and imaging sensors. Therefore, with a DayStar fi lter you can take advantage of the superior resolution and longer focal length of a large telescope for hardly any extra money than you would spend on a fi lter for a small telescope. Indeed, many amateur astronomers have successfully used DayStar fi lters on Schmidt-Cassegrain telescopes from 8 to 14 in. in aperture.

DayStar produce fi lters with a variety of passbands, from 0.8 down to 0.3 Å. All show disc features as well as prominences, although those with the narrowest pass-bands give the greatest contrast on disc features, with the disadvantage that promi-nences are fainter and so require longer exposures to be imaged. The very best (and most pricey) DayStar fi lters, now known as the “Quantum” range, are electrically heated by a small oven controlled by a thermostat. The exact wavelength transmitted by a narrowband DayStar fi lter is strongly in fl uenced by its temperature. An LCD readout on the Quantum fi lters shows the exact wavelength being transmitted down to 0.1 Å, enabling very precise tuning. This precision temperature tuning means that these systems give very high-contrast images, but they are also inconvenient, as they require a power supply. If you have a permanently mounted telescope, you can use mains power (the fi lters come with a transformer enabling you to use the fi lter with your local voltage if you live outside the United States), or if your telescope is portable you can use batteries of the sort used to power telescope drives in the fi eld. The lowest prices are for their “Standard Edition” range of Quantum fi lters. DayStar also offer a “Professional Edition” range in the same series of passbands, intended for professional observatories and so are manufactured to even higher precision than the standard range. (These categories replace DayStar’s former “ATM” [Amateur Telescope Maker] and “University” ranges of premium fi lters.) Professional Edition fi lters sell for nearly twice the price of the Standard Edition.

From the 1980s up until 2010, DayStar also did a line of sub-angstrom fi lters known as “T-Scanners” (Fig. 6.9 ), which adjusted the wavelength transmitted by tilting the etalon using a thumbscrew, rather like the “T-Max” tuning system on Coronado SolarMax fi lters. Because they did not require electrical power, T-Scanners were much more convenient and could easily be used on portable tele-scopes. However, in 2011 DayStar replaced the T-Scanners with the new “Ion” range of fi lters, which are electrically controlled, like the more expensive “Quantum” series. Thus some of the convenience and portability of the old T-Max fi lters was lost, although the power supply is 12 V DC, so Ion fi lters can be powered by batteries of the sort used to power telescope drives in the fi eld. Ion fi lters range in price. It may also be worth looking out for a T-Scanner on the second-hand market, as this line of fi lters was discontinued only recently, and before the advent of Coronado and other manufacturers they were by far the most popular sub-angstrom fi lters on the amateur market – though be aware that some second-hand fi lters may have outlived their manufacturer’s warranty!

118 6 Observing the Chromosphere

Fig. 6.9 A DayStar T-Scanner sub-angstrom H-alpha fi lter, a popular model with amateur astronomers for many years and still used by many solar observers. ( a ) The fi lter unit housing the etalon. ( b ) The fi lter unit mounted on a Meade SCT (Photographs by Richard Bailey)

If you want to increase the contrast of fi laments and other disc features in your H-alpha image but cannot afford the extra etalon required for double-stacking, there is a very cheap method: simply thread a good-quality Moon fi lter into the front of your eyepiece (Fig. 6.10 ). This greatly reduces the glare seen in the H-alpha image, especially when the sky is very transparent, the image is bright and the increase in contrast is very noticeable. Plages and active region fi laments within sunspot groups stand out particularly well with a Moon fi lter. Of course, this “poor person’s” alter-native to double-stacking does not reduce the fi lter’s passband as double-stacking

119Equipment for Observing the Chromosphere

Fig. 6.10 Attaching a Moon fi lter in front of the eyepiece will increase the contrast of disc features in an H-alpha image – though it does not narrow the passband and the prominences will be fainter (Photograph by the author)

does, so you won’t see the same level of chromospheric detail as you would with an extra etalon, but it is surprising how much extra detail you can see with this method. Prominences are also fainter when you use a Moon fi lter – but then, dou-ble-stacking also reduces the brightness of prominences. The same effect can be seen when thin clouds pass over the Sun: the image is dimmed and plages and fi ne fi lament detail temporarily show up with better contrast.

If you are content to view just prominences, and want spectacularly detailed views of them with a small telescope, the “H-alpha Coronagraph” manufactured by Baader Planetarium (Fig. 6.11 ) may be worth considering. Like the DayStar T-Scanner fi lters, the Baader coronagraphs have now been discontinued but may occasionally be found on the second-hand market and are so useful that they are worth looking out for. This instrument screws into an ordinary refractor and works on the principle of the coronagraph designed by Lyot, using a specially made metal disc to produce an arti fi cial eclipse of the Sun. But although called a coronagraph, the Baader instrument shows only prominences, not the corona. At its heart is an H-alpha fi lter with a 1.5 Å passband – too wide for most disc features to be visible but more than narrow enough to reveal prominences.

Unlike other H-alpha fi lters, when used with telescopes of 80 mm aperture or smaller the Baader coronagraph does not require an ERF. Instead, the system’s effective f/ratio is increased by special lenses inside the device. The beauty of this is that the telescope can be used at full aperture, allowing prominences to be seen

120 6 Observing the Chromosphere

in exquisite detail, the resolution limited only by the aperture of the telescope and the seeing conditions. Because the Sun’s surface is masked by the occulting disc, the prominences appear extremely bright. The effect really is like a total solar eclipse – although without the corona, of course. The coronagraph can be used with an ordinary 1¼ in. star diagonal and eyepiece, and it also contains a standard thread for a T-ring, enabling any DSLR (or 35 mm) camera to be attached for photograph-ing the prominences.

The system does have some disadvantages. First, it is less compact than other fi lter types, the version for an 80 mm refractor adding some 200 mm to the length of your tube. It can also be inconvenient in that it consists of three parts, which need to be threaded together before being screwed onto the telescope – a time-consuming process if your observing time is limited. Because the Sun’s apparent diameter varies slightly during the year, the coronagraph comes with a set of six differently sized occulting discs. A different disc has to be threaded into the device every 2 months, to account for the Sun’s changing diameter. Finally, to keep the Sun hidden behind the disc you need a driven equatorial mount that is precisely aligned with the pole. Accurate polar alignment is dif fi cult during the daytime, when you cannot see the stars.

This device is ideal for an observatory, where the telescope is permanently set up and polar aligned. But these disadvantages are outweighed by the incredibly detailed and high-contrast views of prominences it can give. The Baader corona-graph is not suitable for imaging with a webcam or other high-resolution video

Fig. 6.11 A Baader H-alpha coronagraph on an 80 mm (3.1 in.) refractor (Photograph by the author)

121Equipment for Observing the Chromosphere

imaging camera, because these devices have too small a fi eld of view. This may be one reason why this instrument has been discontinued. Another may be the fact that other, cheaper fi lters that have come onto the market allow prominences and also disc features to be viewed much more easily. But unless you are prepared to pay for a very large fi lter, to get high-resolution images of prominences surrounding the whole solar disc with a small telescope, the Baader device still leads the way. The device excels when used with DSLR cameras; the image of the giant prominence shown in Fig. 6.2 was taken with a Canon 300D DSLR attached to a Baader corona-graph on an 80 mm refractor.

Choosing an H-Alpha System

What H-alpha fi lter or telescope you should buy depends on your interests in H-alpha work. An essential point to bear in mind with H-alpha fi lters is that because they are specialist items, many of them are not mass-produced. Popular, entry-level systems such as the Coronado PST are available from stock at many dealers, but some more advanced models have waiting lists ranging from a few weeks to several months.

If you are on a restricted budget, or are only casually interested in H-alpha observing, the Coronado PST and Lunt 35 mm are clear fi rst choices. But even if money is no object, the right choice of H-alpha system requires careful thought, because a number of factors need to be taken into account. First, while the views to be had through a very large H-alpha telescope (or a large fi lter mounted on a big telescope) might be mouthwatering, you need to consider how portable such a setup would be. If you have the luxury of an observatory, portability is not an issue, but you still need to consider whether you have enough space for a big H-alpha system as well as your existing telescope.

If you need to set up your telescopes every time, you will probably want to observe the Sun in white light as well as H-alpha, and it would be very inconvenient to have to set up two telescopes each time you want to observe the Sun. A sensible, and very popular option, is to piggyback your H-alpha telescope on your white-light one – but is your mount strong enough to accommodate two telescopes, espe-cially if one of them is of a substantial size? Even if you have a large budget and are a serious solar observer, it might be better to buy a compact H-alpha telescope, say a Solarview 50 or a Coronado SolarMax 60, and mount it piggyback on your white-light telescope. Alternatively, it might be worth considering buying an ordinary small refractor and mounting an H-alpha fi lter onto it, then piggybacking this arrangement on your white-light telescope. This way you can easily compare the white light and H-alpha views without the inconvenience of changing fi lters. It is also good value for money, in that you have two telescopes for the price of one: at night you can attach a DSLR camera to a wide-field refractor and use it (with H-alpha fi lter removed) for photography of comets and deep-sky objects. You can-not do this with a dedicated H-alpha telescope!

122 6 Observing the Chromosphere

As well as the issue of portability, another disadvantage with a big H-alpha telescope or fi lter arrangement is that in many climates an aperture greater than 100 mm can seldom be used to its full advantage for solar observing, at least as far as ordinary visual observing is concerned, because seeing during the day rarely permits detail smaller than 1 arc sec to be resolved. An arc second is the approxi-mate resolution limit of a 4-in. telescope, so you will only rarely get the bene fi t of an instrument larger than this. However, nowadays webcams and other, more spe-cialized astronomical video cameras can sometimes capture fl eeting moments of good seeing that elude the visual observer, so it can be worthwhile using a larger instrument for this type of imaging.

If you are really serious about high-resolution imaging, you may well have a permanently mounted telescope setup anyway, so portability will be less of a problem. In short, if high-resolution imaging is your main interest, you may wish to consider a large telescope or fi lter. But if your interest is primarily in visual observing, and/or you use portable equipment that needs to be set up and taken down each time, a small H-alpha instrument is the best choice. If you have money to burn but don’t have a permanent housing for your telescope, invest in the highest possible quality of fi lter in a small aperture.

If you use an H-alpha fi lter on an existing telescope, it is very important that you use the correct mount to attach the etalon to the front of the objective lens. It is certainly not good enough to use a home-made arrangement with cardboard and sticky tape, as you might for a Mylar-type white-light fi lter. H-alpha etalons are much heavier and bulkier than white-light fi lters and will easily fall off the telescope if incorrectly mounted, with disastrous results for the eyesight. The fi lter could also be damaged beyond repair if it falls onto hard ground. In any case, as described above, for an H-alpha etalon to transmit the correct wave-length, it needs to be accurately aligned, and for this a precision-engineered mounting is essential. The fi lter needs to be threaded into an adapter made for your particular telescope, and the adapter in turn needs to be threaded or solidly clamped onto the telescope. You can generally order a suitable adapter when you order the fi lter. Adapters are usually available for most telescopes manufactured in recent years, though if your telescope is really ancient or obscure you might need to have an adapter specially made. However, adapters can be surprisingly costly.

If you buy a fi lter with a front-mounted etalon – i.e., one from Coronado, Lunt or Solarscope – it is also important to order the correct blocking fi lter. Solarscope does blocking fi lters of 15 and 30 mm diameters for their systems. For Coronado and Lunt fi lters, they are available in several sizes from 5 mm to about 30 mm, with the smallest naturally being the cheapest. However, the longer the focal length of your telescope, the larger the blocking fi lter you will need to see the whole Sun at once in the fi eld of view. The image of the Sun in millimeters at the prime focus of your telescope is approximately equal to:

110

f

123Equipment for Observing the Chromosphere

where f is your telescope’s focal length in millimeters. So for a telescope of just 500 mm focal length (such as a small apochromatic refractor), the size of the prime focus image will be:

500

110

4.55 mm=

This means that you can get away with a 5 mm blocking fi lter. But a telescope of 900 mm focal length forms an image 8.2 mm across, making it necessary to use a 10 mm blocking fi lter to see the whole Sun in the fi eld.

A rough rule of thumb is that your blocking fi lter needs to be 1 mm in diameter for every 100 mm of your telescope’s focal length. Using a low-power or wide- fi eld eyepiece makes no difference if you have too small a blocking fi lter: the eyepiece’s own fi eld of view will be heavily cut off, and you will see just part of the Sun through a tiny hole in the center of the fi eld. What matters is the size of the prime focus image of the Sun when it passes through the blocking fi lter. The same is true if you place a camera (with lens removed) behind the fi lter for prime focus photography: only part of the Sun’s image will show up on the picture if the blocking fi lter is too small. For the Sun to fi ll the full frame of a large CCD chip, it may be necessary to purchase a 30 mm blocking fi lter.

Another important factor to consider when ordering an H-alpha fi lter system for use with an existing telescope is whether your telescope will reach focus with the blocking fi lter. As noted above, the star diagonal in which many blocking fi lters are mounted takes up a lot of focus travel, and on some telescopes the Sun’s image may not reach focus, even with the focuser fully racked in. This is something that you might only fi nd out by trial and error, and so it might be a good idea to ask the sup-plier to test the fi lter you want to buy on your model of telescope – or perhaps take your telescope to the supplier and ask if you can test it on the Sun while on the premises. If your telescope and fi lter combination does not reach focus, you may need to use some sort of negative lens, such as a Barlow, to push your telescope’s focal point further out – but here you have to be careful. A Barlow lens effectively increases the focal length, and you may fi nd that the ampli fi ed solar image requires a larger blocking fi lter for the whole image to be visible.

For the same reason, the magni fi cation of the resulting image may be inconve-niently high. I encountered just such a focusing problem when I fi rst mounted the Coronado SolarMax 40 fi lter on my 60 mm Takahashi refractor: the image did not come to a focus, even with the focuser fully racked in. I solved it by placing a 2.5× Tele Vue “Powermate” in front of the blocking fi lter assembly. Because the telescope’s original focal length is just 355 mm, the magni fi cation is only about 35× even with the Powermate. Alternatively, both Coronado and Lunt offer “straight through” versions of their blocking fi lters, which dispense with a star diagonal altogether and so require less focus travel. Solarscope systems come with “straight through” blocking fi lters as standard. As noted above, “straight through” blocking fi lters are the ideal type for photography.

124 6 Observing the Chromosphere

Calcium-K and Other Systems

As noted in the fi rst section of this chapter, Calcium-K (Ca-K) fi lters are specialized instruments, primarily aimed at imagers rather than visual observers, because at 393.4 nm the Ca-K line is at the very edge of the visible spectrum and is beyond many people’s visual range. However, Coronado, Lunt and DayStar have all made systems for observing the Sun in this wavelength and a brief survey of the market is needed if you wish to expand your solar studies into this specialized fi eld.

For several years, Coronado manufactured a Ca-K version of their 40 mm PST, now discontinued. This sold for about the same price as the H-alpha version, and even its physical appearance was almost identical. Indeed, some amateurs have mounted their H-alpha and Ca-K PST telescopes together on the same mount, per-haps with a white-light telescope as well, so that the different views can be directly compared. The 40 mm aperture of the Ca-K PST limited its photographic resolu-tion, and as a visual instrument it was, of course, limited by the violet sensitivity of the observer’s eye and the altitude of the Sun. Also discontinued is Coronado’s much more expensive ($2,999) SolarMax Ca-K 70, a dedicated 70 mm refractor for Ca-K observing.

Lunt does a range of Ca-K fi lters, priced from 845 € to 2,239 €, and some com-plete Ca-K telescopes of 60 mm aperture (1,035–1,935 €). The fi lters and telescopes have passbands of about 2.4 Å. This is much wider than H-alpha fi lters capable of revealing disc details on the Sun, because the Ca-K line in the solar spectrum is itself much wider and so the passband of a calcium fi lter does not have to be as narrow in order to isolate the Ca-K line and provide images that show calcium features.

DayStar pioneered calcium solar fi lters for amateur astronomers in the late twentieth century, when they developed Ca-K fi lters that complemented their T-Scanner range of H-alpha fi lters. Now they manufacture two types of calcium fi lter. The most expensive ($5,800) is the Quantum Ca-K 2.0 Å passband fi lter, whose wavelength transmission is electronically controlled and can be viewed on an LCD readout, as on DayStar’s premium Quantum H-alpha fi lters. As with the latter fi lters, the Quantum Ca-K model is best suited to a permanently-mounted telescope. DayStar also now offers a calcium fi lter that is more suited to portable instruments, because it requires no power. However, this model ($4,000) transmits not Ca-K but the Ca-H line, not far from Ca-K in the solar spectrum at 3,969 Å.

Because Ca-H is slightly further from the violet end of the visible spectrum, it should be easier to see the Sun visually in this wavelength. Calcium fi lters do not need to be used at focal ratios of f/30 or longer; stopping down to f/15 or f/20 produces a beam that is suf fi ciently parallel. Nor do they require an ERF over the front of the telescope – indeed, the red ERFs used for H-alpha observing will not work as they do not transmit calcium light. However, you do need to use a UV/IR cut fi lter that is attached to the calcium fi lter just in front of the etalon. This type of fi lter re fl ects ultraviolet and infrared light back out of the front of the telescope, and for this reason is sometimes called a “hot mirror.” Because there is no ERF in front

125Prominences and Filaments

of the telescope aperture, calcium fi lters cannot be used with SCT telescopes or any telescope that has optics inside the tube, in front of the “hot mirror” (this includes some premium apochromat refractors).

Baader Planetarium produces a calcium “K-Line” fi lter, but this is very different from the other fi lters described here. Whereas the Coronado, Lunt and DayStar calcium fi lters have passbands between 2 and 5 Å wide, the Baader fi lter has a passband of 80 Å. This means that it transmits far too much of the Sun’s light for safe observation and must be used either with a full aperture Baader AstroSolar fi lter (of the lighter density used for photography – see Chap. 7 ) or with a Herschel Wedge, which removes much of the Sun’s light and heat (see Chap. 2 ). Even then, the Baader Ca-K fi lter is not safe for visual observation, as it transmits too much ultraviolet light. Like other Ca-K fi lters, the Baader fi lter is intended for photogra-phy, not visual observing. Its performance may not match that of the specialized top-of-the-line Ca-K fi lters, but at just a few hundred dollars it may be a worthwhile addition to your equipment arsenal if you are curious about what is happening on the Sun in Ca-K but are not sure whether to commit thousands of dollars in buying a high-end calcium fi lter.

DayStar also offer sub-angstrom fi lters transmitting one of the sodium D lines at 5,895.9 Å in the yellow part of the spectrum. This wavelength is especially useful for studying fl ares. DayStar sodium fi lters come in two versions: a T-Scanner version at $4,500, which requires no power (similar to the traditional T-Scanner H-alpha fi lters), and an electronically controlled “University” version for $7,200. Both ver-sions have a very narrow 0.4 Å passband, which is essential for isolating the very fi ne sodium line. DayStar also makes a fi lter wheel that allows you to load up to four monochromatic fi lters and then immediately compare the views in different wave-lengths by switching fi lters. But its fi ve- fi gure price tag puts it beyond the range of all but professional or public observatories or the most af fl uent amateurs.

Prominences and Filaments

Prominences can be divided into two broad categories – quiescent and active . Quiescent prominences are long-lived, lasting up to several months and surviving several solar rotations. They can be very large, running to hundreds of thousands of kilometers in length, although are usually only a few tens of thousands of kilometers high.

The commonest type of quiescent prominence is known as a “hedgerow” and, as can be seen from the image in Fig. 6.12 , this is not a bad description, as it is long and low, connected to the Sun by a few short ‘trunks.’ Hedgerows can occur any-where on the Sun, including very high latitudes where sunspots never form. Over time the Sun’s differential rotation tilts them so that they assume an east–west orientation and sometimes a number of such east–west tilted prominences join together to form a “polar crown,” visible on the disc as a huge, long string of fi laments across the Sun. Filament numbers wax and wane in a cycle, but they are out of step with the sunspot cycle, with their numbers being highest shortly after

126 6 Observing the Chromosphere

sunspot minimum. This is one feature that makes H-alpha observing so appealing: even at minimum, when there are few or no sunspots visible, there are usually some prominences.

Quiescent prominences are very stable, changing their appearance only slowly during their long lives, but they end their lives spectacularly when they lift off the Sun at speeds of up to several hundred kilometers per second and fi nally break up, all within the space of a few hours. When this happens, astronomers call the promi-nence an eruptive prominence or a prominence eruption . Astronomers still do not know what causes this type of prominence to suddenly erupt after several months of stability, but they believe that some eruptions are triggered by fl ares elsewhere on the Sun. Many amateur astronomers have witnessed prominence eruptions using H-alpha fi lters, and they are very satisfying to record with a series of sketches or photographs made at intervals of, say, 15 min.

A second type of prominence is associated with sunspots and appears as a loop above the spots. This is technically known as an active region fi lament , as when seen on the disc it appears as a fi lament winding through the sunspot group. Active region fi laments are generally smaller and darker than quiescent fi laments (Fig. 6.13 ). Astronomers believe that such a fi lament marks the line separating opposite magnetic polarities in a sunspot group. Active region fi laments can persist for days or weeks after the sunspots have disappeared and may evolve further to become quiescent prominences.

Active prominences are associated with solar fl ares and can change their appear-ance over periods of a few minutes. A common type is a surge , which appears as a

Fig. 6.12 A typical hedgerow prominence (Image by Dave Tyler)

127Prominences and Filaments

short-lived “spike” at the solar limb. This involves material shooting up into the (invisible) inner corona in a straight path at speeds similar to those in erupting quiescent prominences and then descending back into the chromosphere in the same path. The fastest-changing prominences of all are known as sprays , which indeed resemble a spray of glowing red hydrogen, reminiscent of a volcanic erup-tion. The contents of a spray prominence can move outwards faster than the Sun’s escape velocity and so are ejected from the Sun altogether. Flares can also produce small, brilliant loop prominences, composed of material descending back into the chromosphere. Note that despite their spectacular eruption speeds, prominences do not generally change their appearance before your very eye, due to the Sun’s great distance from us. It usually takes a minute or two for changes to become apparent in a telescope.

At fi rst glance, the chromosphere appears as a thin, smooth band of light around the limb of the Sun. However, if you examine it carefully using an H-alpha setup with reasonably good resolution you may notice that its outer edge appears some-what furry. It is, in fact, full of tiny, prominence-like spikes, known to astronomers as spicules . These behave like prominences in that they are composed of material rising in columns and then gradually fading from view, but they are part of the general structure of the chromosphere and form all over the Sun. They are not associated with sunspot groups or active regions and so are not counted as promi-nences in prominence statistics.

Fig. 6.13 H-alpha image of a sunspot group, showing an active region fi lament. Note that this type of fi lament is smaller and darker than typical quiescent fi laments (Image by Dave Tyler)

128 6 Observing the Chromosphere

Counting Prominences

The simplest form of useful observation you can make of prominences is to count the number of prominences you see each day and work out a prominence MDF at the end of each month, just as you would do for sunspot groups in white light. The exact method of counting varies with different observing organizations. The BAA Solar Section method is to divide the limb into latitude zones of 5°, each zone counting as one prominence active area if it contains one or more prominences. It can, however, be dif fi cult to establish what is 5° in your eyepiece, unless you have a reticule or some other way of making positional measurements (see below).

Whatever method you choose, prominences do tend to occur a good deal further apart than this, whether they form singly or in clusters of several prominences, and it is usually quite easy to tell whether or not a prominence counts as an active area. Occasionally you will see prominences that appear to be detached altogether from the chromosphere, and these do not count as active areas. Neither should very small prominences less than 30 arc sec high be included in prominence counts. Deciding whether a prominence is tall enough to be countable is not always easy, unless you have a micrometer or reticule for precise measurement, but it is some-thing that can be learned with experience. Try to imagine what a planet 30 arc sec across, such as Venus in its crescent phase, looks like at the image scale in your H-alpha instrument. If you are unsure as to whether a prominence is tall enough to be countable, or whether a detached prominence is, in fact, very faintly con-nected to the chromosphere, it is best to give it the bene fi t of the doubt and count it, as it is still an indicator of activity and this is what we are trying to measure. You should always make a note in your observing record when you have made such a decision.

Prominence Position Measurements

Drawings and position measurements for prominences are much harder to do than for sunspots. In H-alpha, we do not have the convenience of projecting the Sun’s image onto a grid and copying down the positions of features. The images pro-duced by an H-alpha fi lter are much too faint to project and so we are constrained to drawing what we see through the eyepiece. One method is to draw the promi-nences after you have drawn the sunspots using the white-light projection method. You can then use the positions of the spots as a reference frame on which you can plot prominence positions. But accuracy is always a compromise with this method, especially if there are not many sunspots present.

An ideal measuring instrument is a reticule, either silhouetted against the Sun’s disc or an illuminated reticule seen against the black sky surrounding the Sun, with angles marked out at intervals of 5° or less. Using such a reticule you could draw prominences in the same way as you would draw sunspots in white light.

129Prominences and Filaments

First, establish the orientation of the image using the drift method. Take care when establishing the directions of north and east: if you are viewing the H-alpha image through a star diagonal (as found in many Coronado and Lunt systems), the orienta-tion will be different from that found when observing straight through or when viewing a white-light projected image. The orientation also changes with the time of day and with the angle at which the star diagonal is pointed (see Chap. 3 ). Then, using a protractor, draw a pattern identical to that of the reticule around the edge of the grid you use for solar drawing, and overlay a sheet of thin paper as before. You can now copy down the positions of the prominences seen in the eyepiece as easily – and as accurately – as you would for sunspots on a projected image. The trouble is that a suitable reticule can be hard to come by.

One excellent device is an eyepiece designed for guiding long exposures of deep-sky objects. This has an illuminated reticule that includes, among other things, a 360° circle around the circumference of the fi eld, marked off at 5° intervals – ideal for prominence drawing. Eyepieces of this type containing an illuminated circular reticule and crosswires, are available from Baader Planetarium, Meade and Celestron. Unfortunately, the main purpose of these eyepieces means that they give a relatively high magni fi cation, their focal lengths usually being 12.5 mm or shorter. This is no problem on a small telescope with a short focal length; for example, the Coronado PST instrument has a focal length of only 400 mm, which gives a magni fi cation of only 32× with a 12.5 mm eyepiece. But in instruments with longer focal lengths the magni fi cation will be too high to show the whole Sun in the fi eld of view – essential for measuring prominence positions. Another possibil-ity is the pole- fi nding eyepieces supplied with some up-market fi nderscopes, as these tend to have lower magni fi cations. You could also try your luck at government surplus optical suppliers, which sometimes sell reticule eyepieces originally intended for use with microscopes or military optics. Other amateurs have invented various ingenious ways of measuring prominence positions indirectly, but you may fi nd that the easiest way of measuring prominence positions is from digital images. Information about correctly orienting a digital image is given in Chap. 7 .

A somewhat easier way of recording prominences is to make a detailed draw-ing of a particularly interesting prominence. Because we are not measuring their positions relative to the whole Sun, precise measurements with a reticule or other means are not required. Establishing the directions of north and east using the drift method is all that is necessary. We can use relatively high magni fi cations here, because we are not viewing the whole Sun. While drawing by hand is not as accurate as imaging, it is often more convenient, as setting up the camera and focusing the image can be very time-consuming, particularly if clouds are threat-ening to roll in and end observations, or the prominence under observation is changing very quickly. Making an accurate prominence drawing is not easy, how-ever, because when viewed through a good instrument prominences can be very intricate. It is best to start by drawing relatively small and simple prominences, before progressing to more complex types. To ensure consistency and ready com-parison with other drawings, try to do all your prominence drawings to the same scale. Represent the Sun’s limb by drawing a segment of a circle of known diameter.

130 6 Observing the Chromosphere

To begin with, you could use a 6 in. protractor for this and draw prominences to the same scale as full-disc drawings in white light. Begin a drawing by outlining the basic shape of the prominence as accurately as possible and then fi ll in the fi ner details.

It is possible to estimate the heights of prominences, using drawings or images. Simply measure the height of the prominence on your image in millimeters or inches and then divide it by the diameter of the disc you are using. Multiplying the result by the Sun’s true diameter (1,392,000 km, or 864,000 miles) gives the actual height of the prominence. A simple measurement such as this reminds us of the vast scale of solar features. Even quite small prominences dwarf Earth, whose diameter is 12,756 km. Using this method, the huge prominence in Fig. 6.2 can be estimated to be 80,000 miles high and 210,000 miles long – nearly the distance from Earth to the Moon!

If you have a fi lter capable of revealing disc features you can record fi laments visually by drawing them on a white light sunspot drawing form, after you have made a sunspot drawing and so using the spots as a guide to positioning the fi laments. It is interesting to distinguish different types of fi laments and their equivalent appearance as prominences on the limb. Quiescent fi laments tend to be long and thick, although (as noted above) they are fainter than the active region types. When they are near the limb, fi laments can appear three-dimensional and quiescent fi laments sometimes show the “hedgerow” appearance of quiescent prominences. The dark, thin active region fi laments often have a winding, snake-like appearance, and they sometimes show up as loop prominences at the limb. Note also how the chromospheric structure changes close to sunspot groups. The granular structure is distorted and often shows the familiar pattern formed by iron fi lings over a magnet.

One exciting aspect of observing fi laments is watching them rise off the surface of the Sun. Like prominences, fi laments move too slowly for their motion to be seen directly, but you can tell they are moving using the Doppler effect – the compres-sion or lengthening of the wavelength of light emitted by an object moving towards or receding from us. In astronomy, the Doppler effect is best known for its effect on the light of distant galaxies receding from us at very high velocity. As the wave-length of light emitted by the receding galaxies lengthens, the light becomes more and more red, because red light is simply light with a longer wavelength. Thus the effect is known as “redshift.” Conversely, light emitted by objects moving towards us is blue -shifted, as its wavelength is reduced.

We can observe blue-shifting on a smaller scale on the Sun. When an H-alpha fi lter is tilted slightly away from the focal plane, its peak transmission is shifted away from the H-alpha line, towards the blue end of the spectrum. If a fi lament is rising towards us, it appears fainter when the fi lter is aligned square-on to the opti-cal axis but becomes more prominent when the fi lter is tilted. The effect is most dramatic when a prominence erupts while it is visible as a fi lament on the disc. The fi lament then rapidly disappears – not because it has broken up and its con-tents dispersed but because it is moving so fast that it has been blue-shifted into the violet part of the spectrum, way outside the range of an H-alpha fi lter. Such an

131Flares

eruption when seen on the disc is sometimes known by the French term disparition brusque – literally, a “sudden disappearance.” Doppler shifting is most noticeable in fi laments near the center of the disc, which are moving directly towards us. Some of the latest H-alpha systems can be tuned to either side of the H-alpha line, and with these it may be possible to see red-shifted material falling back towards the solar surface.

Flares

Flares usually occur in or near sunspot groups (Fig. 6.14 ), but they are occasion-ally seen in active regions before any sunspots have formed or after they have all disappeared. As noted in Chap. 3 , the extremely rare white-light fl ares are most likely to occur in complex groups, of types D, E and F in the McIntosh system. The same is true for fl ares in H-alpha: the strongest and most frequent fl ares are found in D, E and F groups. Very large and complex groups can produce many fl ares during the course of their 2-week transit across the Sun, and if such a group is present it is worth checking it for fl ares with your H-alpha fi lter several times a day, if you can.

In both their frequency and their intensity solar fl ares broadly follow the rise and fall of the solar cycle, although professional astronomers have sometimes observed a slight resurgence in fl are activity in the late stages of the cycle, when sunspot numbers are dropping. At solar maximum, several hundred fl ares can occur each month, although you will not record anywhere near this number with your telescope, because fl ares tend only to last for a few minutes and many will take place when the Sun is not visible. At minimum, only a small handful of fl ares occur in a month, and days or even weeks can go by without a single fl are being recorded. Also, fl ares at this time are generally much smaller than those observed at maximum.

A fl are often begins as a pair of bright points of light in or near a sunspot group. These then brighten rapidly and increase in size, often elongating into short streaks or “ribbons,” before gradually fading from view. Flares generally take much longer to fade out than they do to reach maximum brightness. The majority of fl ares have lifetimes measured in minutes, but very large and active ones can last for several hours. Remember not to confuse the bright plages around sunspots with fl ares; these tend to cover a much larger area and change much more slowly.

Some of the most spectacular fl ares occur at the Sun’s limb (Fig. 6.15 ). Such a limb fl are takes the form of a small but strikingly bright prominence that often comes into view in the space of just a few minutes – indeed, you might not have noticed anything in its position when you fi rst looked at that part of the Sun a short while before. A limb fl are can grow and change its shape dramatically over a simi-larly short period of time. Loop formations sometimes occur and material is occa-sionally ejected from the Sun. Often, though, a limb fl are will simply fade or the contents gradually disperse.

132 6 Observing the Chromosphere

Fig. 6.15 A spectacular limb fl are (Imaged by Dave Tyler on March 8, 2011)

Fig. 6.14 H-alpha image of a fl are in the large sunspot group AR11158, taken by Dave Tyler on February 17, 2011. The fl are is the intensely bright region near the center of the group

133Flares

When you see a fl are, it is important to record several features. Because fl ares are transient events, it is essential to time their appearance accurately, to the nearest minute. Note down the time you fi rst noticed the brightening, the time at which it reached its maximum intensity and the time at which it had faded from view and you could no longer see it. As with all solar observations, make a note of the observing conditions and the telescope and H-alpha fi lter you are using.

Also important is an estimate of the fl are’s position. The most practical way of doing this is to sketch the fl are’s appearance relative to its parent sunspot group. An image is, of course, more accurate, but setting up the camera may take too much time in a short-lived event such as this. When you have made a sketch of the fl are it should then be quite easy to get an accurate estimate of its latitude and longitude from a disc drawing or photograph showing the group’s position on the Sun. If you are able to take correctly oriented, whole-disc images in H-alpha, it should not be dif fi cult to measure the fl are’s heliographic coordinates using one of the online solar position measurement programs mentioned in Chap. 4 .

Astronomers use a system for classifying the appearance of fl ares in H-alpha light based on the area of the Sun’s disc a fl are covers and the fl are’s peak inten-sity. The area in square degrees covered by the fl are when at its peak intensity, sometimes known as the fl are’s “importance,” is described by the numbers 1–4, as follows:

Importance Area covered (square degrees)

1 2–5 2 5–12.5 3 12.5–25 4 over 25

Flares covering an area of less than 2 square degrees are known as sub- fl ares and are denoted by the letter S. The fl are’s intensity relative to its surroundings is assigned using one of three letters: f (faint), n (normal) and b (bright). A brighten-ing not much greater than the surrounding chromosphere would be described as f , while b refers to something strikingly brilliant.

When a fl are is classi fi ed, the number representing its importance is written before its intensity. For example, a large and bright fl are covering 15 square degrees would be class 3 b . In practice, it is dif fi cult for a visual observer to make an accu-rate estimate of a fl are’s intensity and especially its area. It is best to make sketches as described above and leave the fi nal classi fi cation to the director of your observ-ing group. You can, however, make a rough estimate of a fl are’s class using your own sketches and its size relative to the disc drawing.

You may notice astronomical reports that classify fl ares with the capital letters B, C, M or X. This is a classi fi cation system used by professional solar scientists; the letters refer to a fl are’s intensity in X-rays and so are not an exact indicator of its visibility in H-alpha telescopes. However, it is worth knowing what the classi fi cations mean, as they are an indication of a fl are’s power and its potential

134 6 Observing the Chromosphere

in fl uence on Earth. If you see a fl are, it is fascinating to compare a fl are’s X-ray classi fi cation with your own observations of it. Conversely, if you see a live online report of a powerful X-ray fl are, it may be worth looking for it with your H-alpha telescope. Brie fl y, fl ares of X-ray classes B and C are small and have little notice-able effect on Earth. M-class fl ares are medium-sized and can cause minor geomag-netic storms. It is X-class fl ares that astronomers really get excited about and which can cause major geomagnetic storms, including auroral displays – although a big fl are is in itself no guarantee that an aurora will be visible from temperate latitudes. Very powerful X-class fl ares can also cause some of the less pleasant terrestrial effects discussed in Chap. 1 , such as power cuts and malfunctions in essential Earth satellites due to excessive doses of radiation. Each of these classes has nine sub-classes, numbered from 1 to 9, with the sub-class written after the main class, so fl ares are numbered from M1 to M9, ×1 to ×9, etc. So a medium-sized M-class fl are might be more precisely classed as M5. Occasionally, though, fl ares are so strong that they are given numbers higher than ×9. The most powerful solar fl are ever recorded was classed as ×20 it occurred on November 4, 2003, when the great sunspot group pictured in Fig. 3.10 (Chap. 3 ) passed around the Sun’s western limb. Had it occurred when it was near the central meridian, and so aligned with Earth, its consequences for our planet could have been serious.

Observing the Sun’s chromosphere is fascinating – and, indeed, is a whole sub-ject in itself. But don’t let it become all-absorbing and replace white-light observ-ing, as the latter is the main “bread and butter” of amateur observing groups and is the area in which most observations are needed. One rule is always to do white-light observing fi rst and only when this is completed go on to H-alpha if time and weather permit. But the views to be had in H-alpha are so spectacular that it is often tempting to break this rule, especially when activity is high!

135L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0_7, © Springer Science+Business Media New York 2012

Imaging the Sun with a Digital

Camera

Chapter 7

When it is done correctly, solar photography can be very rewarding. It is the most accurate way of recording sunspots and other solar features – although, as explained in Chap. 4 , it is more dif fi cult to do accurate positional measurements using photo-graphs than it is using drawings. Nevertheless, photographs of the Sun’s full disc are excellent for showing what was visible on the Sun on a particular day, and they show the appearance of the Sun’s features far more realistically than drawings ever can.

Digital images are great for demonstrating recent solar activity to astronomy clubs. Photography is the only way of getting a pictorial record of the Sun if there is too little time to make a drawing – for example, if clouds are approaching. Whereas a disc drawing can take half an hour or more to complete if the Sun is active, the Sun can be recorded with an exposure in the region of 1/1,000 s! Of course, mounting the camera and focusing the image takes up some time, but even allowing for this photography is often far quicker than drawing. Especially useful are close-up photographs of individual sunspot groups, as drawings showing the details of complex groups are not easy to do accurately. And if you have an H-alpha fi lter or telescope, it is possible to take dramatic shots of prominences and other chromospheric features.

One of the beauties of solar photography is that, as with other forms of solar observing, superb results can be obtained with just a small telescope. Also, unlike most kinds of night-sky photography, exposures for the Sun are so short that your telescope doesn’t need to be motor-driven to counteract Earth’s rotation. This author did many successful early experiments in solar photography using just a basic 60 mm alt-azimuth mounted refractor costing less than £150. Similarly, the type of camera you use does not have to be sophisticated or expensive. An ordinary digital camera, of the point-and-shoot or digital single-lens-re fl ex (DSLR) variety from your local camera store is all that is needed.

136 7 Imaging the Sun with a Digital Camera

Until the 1990s fi lm was the only medium with which amateur astronomers photographed the Sun, and it remained the main medium until the beginning of the twenty- fi rst century. But in recent years digital photography has more or less entirely supplanted fi lm, and this and the following chapters will focus entirely on digital techniques. Digital cameras are readily available off the shelf at camera stores everywhere, and nearly all models can be used at least to some extent for solar photography.

From the viewpoint of the amateur solar observer, digital photography has many advantages. The most important is that it produces instant results. Once you have taken a picture, you can see it immediately on the camera’s LCD monitor. If you don’t like the picture, you can try again with another shot. With a fi lm camera you had to shoot a whole roll of fi lm and have it processed before you could see the results. Because digital photography allows you to view the results immediately, you can learn from your mistakes very quickly.

A second advantage of digital photography is that you do not need a darkroom to process the images or have a laboratory do the work for you. All you need to process, handle and print out the images is a standard personal computer. Because the images are produced on a computer, they are ready for sending by e-mail as attachments or posting on the Internet. A daily solar image taken with your digital camera would be a great addition to an astronomy club website, as would a “movie” of many solar images taken at daily intervals showing the Sun’s rotation.

This capability of producing instant results and sending them by e-mail makes digital photography very useful for reporting interesting or unusual solar events to solar observing organizations. Digital images can be quickly distributed to other members of these groups. If you are ever lucky enough to see a white light fl are, digital images would be especially valuable. If your digital camera and telescope are aimed at the Sun when such a fl are is in progress, and if you e-mail it quickly enough to your observing group, your image might be one of the fi rst pictures of the event available.

Amateur astronomers usually photograph the Sun with two types of digital camera: the ordinary digital camera as found in camera stores, and the “webcam” – a device originally designed for broadcasting videos over the Internet but which can be adapted to take images through a telescope. The advantage of webcams is that they shoot moving pictures. They do this by shooting many still images per second, the result appearing to the eye like a continuously moving picture. For astronomers, this means that they can capture moments of steady seeing and, pro-vided that the seeing conditions are reasonably good, a few frames will come out super-sharp. Once a video of, say, a few hundred frames has been transmitted to the computer, the best frames can be selected and then electronically stacked to produce an even better image.

Webcams are available from computer stores, but versions designed speci fi cally for astronomical imaging are produced by telescope companies, and some fi rms make quite sophisticated – and expensive – cameras based on the webcam principle for taking really high-quality images. Webcams excel at imaging small objects with very high magni fi cation; this, combined with their ability to shoot moving pictures,

137Choosing a Digital Camera

is why some of the best images of the planets, as well as the Sun, ever taken by amateurs have been made with this type of camera.

Webcams have two major drawbacks. One is that they need to be connected to a computer while they are being operated. The other is that the imaging sensor inside a webcam is very small, and so images taken at a small enough scale to show the whole solar disc, or even a large part of it, have a low resolution. Therefore webcams are less useful if you want to image the whole solar disc or large parts of it. For these purposes, a conventional digital camera is much the best choice. The two types of cameras therefore complement each other: ordinary digital cameras are best for general solar photography, while webcams are the camera of choice for high-resolution close-ups of individual sunspots or other solar features.

There is also a third type of digital camera: the CCD camera speci fi cally designed for astronomical imaging through telescopes. When they fi rst appeared in the 1990s, these cameras revolutionized night-sky imaging, because their high sensitivity greatly reduces exposure times, enabling amateurs with moderate-sized telescopes to capture very faint detail in deep-sky objects and to reach limiting magnitudes unheard of in the days of fi lm. Their short exposure times also greatly improved the resolution of solar (and planetary) images.

If you have such a CCD camera, you can use it to take excellent solar images. In particular, a monochrome CCD camera can be very effective in H-alpha photog-raphy, because the images transmitted by sub-angstrom H-alpha fi lters are rela-tively faint and also because monochrome cameras give the best results in H-alpha (see below). But these cameras take only still images, not movies, and for this rea-son they have been superseded by webcams and the more specialized webcam-like astronomical cameras, whose ability to take movies allows them to take images with a resolution far higher even than CCD cameras. If you have an astronomical CCD camera, by all means try it out for imaging the Sun, but if you are thinking of buying a camera speci fi cally for high-resolution solar photography, you would do best choosing a webcam.

In this chapter we will discuss photography with a conventional digital camera, while the book’s fi nal chapter explores techniques for high-resolution imaging. The latter chapter also discusses how to process and enhance images taken with either type of camera.

Choosing a Digital Camera

“Compact” Cameras

If you already have a digital camera, you can probably use it to some extent for solar photography. If your camera is of the “compact” type (Fig. 7.1 ) and not a DSLR, one feature it must have is a built-in LCD monitor both for composing the picture and viewing the resulting image. With such cameras you do not look through the lens but use an entirely separate view fi nder to frame your shot.

138 7 Imaging the Sun with a Digital Camera

This view fi nder is located in a different position from the lens that takes the picture, which is no problem in everyday photography, but if you are shooting through a telescope this prevents you from seeing what the lens is seeing!

If you are choosing a camera speci fi cally for solar photography, there is a huge range of models available in the camera stores, with prices ranging from tens through to thousands of dollars. However, certain features make some cameras much more suitable for solar photography than others. One of them is the camera’s resolution. Like a dedicated CCD camera, a digital camera uses a light-sensitive CCD (charge-coupled device) chip in place of the fi lm. In some cameras the chip is known as a CMOS (complementary metal-oxide semiconductor) sensor. This uses different technology, but for our purposes it is a light-sensitive chip similar to a CCD. This chip – and the resulting picture made from it – is composed of myriads of tiny elements called pixels. The larger the number of pixels on the chip, the higher its resolution, in the same way as the resolution of photographic fi lm is higher if it has fi ne grain. The resolution of a digital camera is measured in mega-pixels – i.e., the number of pixels divided by one million. So if you see a camera advertised as having 6 megapixels, its CCD or CMOS sensor contains six million pixels – more than enough for quality solar imaging. Generally, a good fi rst camera will be in the 6–8 megapixel range, costing around $350 or a little less.

Although digital cameras are often advertised by the size of their imaging sen-sors in megapixels, this is not the most important feature in a digital camera for the

Fig. 7.1 A typical compact digital camera: a 10-megapixel Canon PowerShot A640. This camera’s LCD view fi nder screen unfolds from the main body of the camera and can be adjusted to a great variety of angles – a handy feature when the camera is mounted on a telescope and the screen is at an awkward angle if left in its conventional position

139Choosing a Digital Camera

solar photographer. The number of megapixels only tells us the resolution of the image. The higher the megapixel count, the higher the resolution – i.e., the smaller the size of detail in arc seconds that can be resolved at a given focal length. But even the cheaper digital cameras nowadays, of just 4 or 5 megapixels, have enough resolution to produce a good-quality 4 in. × 6 in. print without the pixels showing. More important is the depth of the image – that is, the amount of data it contains. The greater the depth of the image, the more shades between the extremes of black and white it contains.

Many compact digital cameras store images in a format known as JPEG , which is a standard format that is readable on any computer. A property of JPEG is that it compresses images – i.e., removes some of the data from the image. This reduces the size of the image fi le on your computer in megabytes, and it is the reason why an image from, say, a 10-megapixel camera may only take up 3 MB on your hard drive. This allows your computer or your camera’s memory to store many more images than would be possible if they were not compressed, but at the cost of some loss of image depth. The greater the JPEG compression, the smaller the image fi le, but the less the image depth, regardless of the size of the image in megapixels. Good image depth is especially important in solar photography, where we are dealing with a range of tones, often over a smooth surface. In extreme cases, a heavily compressed image will show the variation in brightness as a series of layers rather than a smooth grada-tion of brightness. This effect is dramatically visible in the whole-disc image of the Sun shown in Fig. 7.2 . Such an image is not only unsightly to look at but important detail is also lost. It is best to choose a camera that stores images in an uncompressed format, such as TIFF or the camera’s own “RAW” format – although the latter will require the camera’s own software for the images to be readable. If the camera you are considering only does JPEG, ask an experienced camera store assistant how much the images are compressed, as image compression varies from camera to camera.

Digital cameras store images on memory cards – small, removable units slotted into the camera. They are the digital equivalent of fi lm, although when a memory card is full you just have to delete some images (having downloaded them to your computer) rather than replace it. The most common type of memory card nowadays is known as a “compact fl ash” card. It may be worth purchasing extra cards to sup-plement the standard card that comes with the camera, as each image takes up a lot of memory. Beware of cameras that say on their LCD display that you have 1,000 or some other such huge number of photos left. These calculations are usually based on a small to medium image size selection, which means heavily compressed JPEG images – with all the disadvantages described above. When shooting the Sun you need to use RAW, TIFF or the largest possible JPEG image size offered by your camera. Such uncompressed or moderately compressed images take up far more memory. Consider getting at least a 4-gigabyte (4 GB) memory card as a spare. As with RAM on personal computers, you can never have too much memory!

A third useful feature is a reasonably large lens, preferably one with a variable focal length – i.e., a zoom lens. Compact cameras do not have removable lenses, so you need to choose the lens when you choose the camera. Your choice of lens is particularly important in solar photography, because you will be pointing the camera

140 7 Imaging the Sun with a Digital Camera

lens into the telescope eyepiece. A small camera lens with a short focal length shows the telescopic fi eld as very small and perhaps surrounded by blurred telescope parts – a phenomenon known as vignetting .

Another problem caused by small lenses is that some of them show only part of the telescopic fi eld, so that even with a low-power eyepiece part of the Sun’s disc is obscured, or at any rate the camera has to be precisely aligned with the eyepiece for the whole Sun to be visible. Such a shot is obviously rather unattractive. A longer lens excludes most of the surroundings and shows more of the telescope’s fi eld of view, allowing more aesthetically pleasing pictures showing just the Sun surrounded by black sky. The lenses on most digital cameras now in production have zoom lenses, which give image scales ranging from wide angle to moderate telephoto. As you “zoom in” and increase the focal length, vignetting should disappear, and you should see most of the fi eld of view. Beware of zooming in too far, for then the magni fi cation becomes too high and the image becomes faint and dif fi cult to focus. Try to choose a camera with a large zoom factor – that is, with a long range of focal lengths. Many compact cameras now come with a zoom factor of 10× or more. However, make sure the zoom factor advertised with the camera refers to “optical zoom” and not “digital zoom.” The latter merely increases the size of the pixels, leading to a grainy picture.

Fig. 7.2 A heavily compressed JPEG image, occupying only about 100 KB on the computer’s hard drive but rendered almost useless by the variation in tone caused by the limb darkening being reduced to a series of layers. The original image was taken on a 6.3-megapixel DSLR, demonstrating that the degree of image compression is more important than the size of the image in megapixels

141Choosing a Digital Camera

When you take a picture of the Sun (or anything else) through your telescope, the act of releasing the shutter will cause a slight movement of the camera, known as “camera shake,” and this can cause a blurring of the image. The effect of vibra-tion on image quality may affect your choice of camera. The problems caused by camera shake depend on how you mount the camera to take pictures, as is described in the separate section below. If you mount the camera on a separate tripod and point it into the eyepiece to take photographs, camera shake should not, in theory, be a problem, because the vibrations caused by your hand pressing the shutter are not transmitted to the telescope. But in fact, it can still cause trouble, especially if (as described above) the camera has to be precisely aligned to get the whole Sun in the picture.

Vibration can cause great problems if you attach the camera to the telescope in any way, because even tiny vibrations are enough to blur the image. Digital SLRs allow the shutter to be released using a cable or electronic remote control, but because they are designed for everyday amateur photography, many compact cam-eras have no such facility and so you have to squeeze the shutter release by hand. Fortunately, digital cameras often have a self-timer facility, designed for allowing the photographer to appear in the picture, which gives an interval of several sec-onds for any vibrations caused by the hand to die down before the shutter is opened. Try to make sure a camera has this feature before buying it.

Another feature worth considering when buying a camera for solar photography is the ability to control the shutter speed manually. Many of the cheaper models offer only automatic exposure, which is fi ne for everyday photography but inade-quate for the Sun, particularly if you are shooting a small whole-disc image sur-rounded by a lot of black sky, as the camera will then overexpose the Sun. Features such as “exposure compensation,” “spot metering” and “center-weighting” found on some cameras may improve the exposure accuracy, but there is no substitute for being able to set the camera to a speci fi c shutter speed.

It is important to check over the camera’s LCD screen, for you will be using this to center and focus the Sun’s image. Try to select a model with as large a screen as possible, as a large screen will have more pixels and therefore a higher resolution and so will enable you to achieve a more accurate focus. The screen should also give a bright image, because otherwise it will be hard to see in full sunlight. Many camera screens have adjustable brightness, and it is worth asking in the camera store if you can try the screen of your selected camera at maximum brightness before you part with your money. Best of all is a screen that can be tilted with respect to the main body of the camera. Not only does this allow the LCD to be tilted away from direct sunlight, but it also lets you adjust it to a convenient viewing position – a very handy feature when the camera is mounted on a telescope, when the position of the screen can be awkward for viewing, particularly if you are shoot-ing when the Sun is high in the sky.

All digital cameras require batteries to function, and solar photography can be problematic in this regard, because whereas in everyday photography you can compose your shot with the camera’s view fi nder, when shooting through a tele-scope you need to use the LCD screen all the time, which consumes a lot of power.

142 7 Imaging the Sun with a Digital Camera

Even a short imaging session can completely drain the batteries. To avoid having to constantly replace batteries, buy a battery re-charger and re-chargeable batteries when you purchase the camera. Some models come with their own high-perfor-mance re-chargeable battery pack, and this is another feature to consider when buying a camera. If your camera is powered by standard AA-size batteries, try to use nickel metal hydride (NiMh) batteries, as they last somewhat longer after each charge than do nickel cadmium (NiCD) cells. It is a good idea to obtain two sets of batteries and keep them fully charged, so that if one set runs out in the middle of a session you can still keep going. Many cameras have a facility for an adapter to allow the camera to run on mains electricity, but this is not recommended when using a camera outdoors, as it often involves trailing a live cable over damp grass and perhaps metal telescope or tripod parts.

Digital SLRs

Compact digital cameras are fi ne as a quick and easy way of photographing the Sun, and if you are just starting out in solar photography or are just casually inter-ested in this fi eld, they will give very satisfactory results. But if you are serious about imaging the Sun, the best type of digital camera is, without doubt, the digital single-lens re fl ex (digital SLR or DSLR) (Fig. 7.3 ). DSLRs were initially aimed at the professional photography market and used to be beyond the price range of most amateurs, but in recent years their price has dropped dramatically. The cheapest DSLRs now sell for around $650 brand new – no more than some high-end com-pact digital cameras. In fact, the best DSLRs for the amateur solar photographer are not the professional-grade cameras at the top end of the market, with four- fi gure price tags and image sensors the same size as a frame of the old 35 mm fi lm (actu-ally 24 mm × 36 mm across) or even larger. Such cameras can be too heavy for use on small telescopes, since they upset the balance of the telescope too much. Also, a large image sensor can actually be a disadvantage, because if it is larger than your telescope’s drawtube you may not use all of the sensor’s surface area and so the image will be vignetted at the edge. Cheaper models, such as a 6.3-megapixel Canon EOS 300D, or a later model such as the 550D, are much lighter and have somewhat smaller image sensors. All the major camera brands – including Canon, Nikon, Olympus and Pentax – now make DSLRs, with Canon and Nikon standing out as the most popular makes among amateur astrophotographers.

DSLRs have several distinct advantages over ordinary digital cameras. The fi rst, and greatest, is that you focus the image using the view fi nder, not the LCD screen. This allows the image to be focused much more accurately and also greatly reduces power consumption, since the power-hungry LCD monitor is not used. You only use the LCD for viewing the resulting image after you have taken the exposure. At a moderate extra cost, you can also attach a “focusing magni fi er” to the view fi nder; this allows even more accurate focusing and also lets you compose your shot from a more comfortable angle.

143Choosing a Digital Camera

Almost as important as through-the-lens (TTL) focusing is the ability to remove the lens altogether. This allows you to attach the camera to the telescope much more solidly than is possible with a compact camera, since you can clamp the camera to the telescope using standard-speci fi cation adapters and rings that have been avail-able from telescope suppliers for many years. These are the same adapters that were used to mount 35 mm SLR cameras to telescopes. The ability to remove the lens makes a DSLR a good value for the money, because if you buy a second camera (of the same make), you need only buy the camera body, as you can use the same lenses with the new camera as you did with the fi rst.

These cameras are also a good investment if you plan to do night-sky imaging as well, as unlike most ordinary digital cameras they have a “bulb” setting for long exposures. Compact cameras generally have a maximum exposure time of a few seconds, which severely limits their usefulness for night-sky photography. Indeed, full manual exposure control is a separate major advantage of DSLRs over compact cameras. They offer a huge range of manual exposures, from “bulb” down to 1/4,000 of a second and sometimes even faster.

DSLRs give you further control over the brightness of the resulting photograph in that they also have “ISO” settings that allow you to vary the image sensor’s sen-sitivity to light. Only some compact cameras have this feature. The ISO setting is equivalent to the sensitivity or “speed” of the fi lm in the pre-digital era. So an ISO setting of 50 or 100 is equivalent to a “slow” fi lm and indicates relatively low sensi-tivity. ISO 400 is a medium speed, while ISO 1600 is equivalent to a “fast” fi lm and

Fig. 7.3 The author’s 6.3-megapixel Canon EOS 300D digital single lens re fl ex (DSLR) cam-era, shown with 18–55 mm zoom lens separated from the camera body. With the lens removed, cameras of this type can be securely clamped to the telescope with suitable adapters to take high-quality images of the Sun (and other astronomical objects)

144 7 Imaging the Sun with a Digital Camera

allows very short exposures. Some cameras have settings of ISO 3200 or even faster, but the faster the ISO setting, the more the amount of electronic “noise” – i.e., unwanted bright and dark pixels – you will see on the image. Noise is especially noticeable if you are shooting something with a smooth, amorphous texture, such as the Sun in white light. For the best solar photographs, it is best to use a slower ISO setting and a less dense fi lter, as is explained below in the section on fi lters.

Similarly, image depth is much less of a problem than it is with compact cam-eras. Most DSLRs can output uncompressed images, either in TIFF format or in the manufacturer’s proprietary RAW format. Even the JPEG options in DSLRs, at least in the larger fi le sizes, suffer from far less compression than do their equivalents in compact cameras, perhaps because DSLRs are aimed at professional and serious amateur photographers, rather than the everyday “snapping” for which many com-pact cameras are designed. In fact, with a Canon 300D, the largest JPEG fi le size is quite satisfactory for all but the most exacting purposes in solar photography, and this is the most convenient option in practice, since RAW fi les have to be converted to TIFF in order to be viewed and processed on a computer. The huge size of the resulting TIFF fi les means that a lot of them can take up a signi fi cant amount of space on the computer’s hard disk. The image sensors in DSLRs tend to have a higher megapixel count than conventional digital cameras. However, as noted above, the amount (or rather the lack of) image compression is more important than the number of pixels. An 8-megapixel camera that produces uncompressed or nearly uncompressed images is far preferable to a 12-megapixel camera that only offers heavily compressed pictures.

Finally, DSLRs allow you to easily release the shutter remotely. Instead of the manual cable release that was formerly used with 35 mm cameras, in DSLRs the shutter is released electronically, using either a “hard-wired” remote release that plugs into a socket on the camera body, or a separate remote control that operates the camera via an infrared beam, like the remote control for a television. The latter type of camera control is more expensive and requires separate batteries, but it allows the camera to be controlled from a much larger distance than a hard-wired release, which is typically no longer than a traditional cable release – about 20 in. DSLRs can also be remotely controlled via your computer, or a laptop if your tele-scope is set up in the fi eld, but this requires a lot more setting up than simply using the camera on its own, and one of the beauties of using a shop-bought digital cam-era is the ability to use it without a computer or external power.

Telescopes and Mounts

You can use any of the three main telescope types for solar photography. In the-ory, an apochromatic refractor should have the edge: it has no central obstruction and so the image contrast should be higher, and it does not have the chromatic aberration inherent in simple refractors. But in practice all telescopes can give excellent results. Unless you are using a Herschel wedge, there are no prob-lems with heat build-up in Newtonians, Schmidt-Cassegrains and Maksutovs.

145Telescopes and Mounts

(As discussed in Chap. 2 , a Herschel wedge should not, of course, be used in catadioptric telescopes, because of the danger of delicate internal telescope parts overheating).

If you use a Newtonian, however, you may fi nd that you cannot rack the focuser in far enough for the image to reach focus in your camera’s view fi nder if you are using a DSLR. One remedy for this is to substitute your existing focuser with a “low-pro fi le” device intended for astrophotography, which enables the camera to be placed closer to the tube. Alternatively, use a Barlow lens between the camera and the focuser. This, of course, increases the size of the image and the effective f/ratio, and therefore the exposure, but as many modern Newtonians are relatively “fast” (f/8 or faster), this is not usually a problem.

Whatever type of telescope you use, if it is larger than 150 mm (6 in.) you need to stop down the aperture. At most observing sites, the seeing conditions during the day rarely permit resolution higher than the theoretical resolving power of a 100 mm (4 in.) telescope – i.e., about 1 second of arc. Solar images taken at full aperture using a large instrument tend to be blurred and have a washed-out appear-ance. Stopping down will improve the contrast and sharpness of the image, even though the theoretical resolution will be lower. If you have a Mylar-type fi lter, you can stop down the aperture by making an off-axis mask from cardboard as described in Chap. 2 . Mount the fi lter by making two layers of cardboard, each with a hole of the same size, and sandwiching the fi lter material between them. Glass fi lters usu-ally come ready-made, with the fi lter already mounted in the mask, which is sized to suit a speci fi c aperture and make of telescope.

More important than the telescope type is its mount. If you are to reduce image-blurring vibrations to an absolute minimum, you must use as steady a mount as possible. This is yet another reason for avoiding very cheap telescopes, as many of them are sold on fl imsy mountings. Fortunately, many “serious” astronomical tele-scopes are nowadays sold on mounts designed with astrophotography in mind, and the requirements for solar photography are not as stringent as those for night-sky work, which demands much longer exposures. My 80 mm Vixen refractor is mounted on Vixen’s well-known Great Polaris mount, which is often used with somewhat larger refractors. It helps to ensure that there are no loose joints in the mount – for example, by tightening the thumbscrews connecting the tripod to the equatorial head.

As noted at the beginning of this chapter, a motor drive is not essential for solar photography. To avoid vibration, you need to stick to exposures of 1/125 of a sec-ond or faster, and these are too short for the image to be blurred by Earth’s rotation, even at very long focal lengths. It has to be said, however, that a drive is very con-venient, as when you are setting up and focusing the camera the Sun quickly disap-pears from view if the telescope is not driven. Electronic slow motion controls also make it easier to center individual sunspot groups at high magni fi cation. A driven telescope with electronic slow motions is a must if you are using a dedicated astro-nomical CCD camera or webcam, as the image sensors on these cameras are usu-ally very small, and it would be extremely dif fi cult to center a sunspot or other target without adequate slow motion controls.

146 7 Imaging the Sun with a Digital Camera

Filters

It is wise to begin with a word of caution here: SOLAR PHOTOGRAPHY THROUGH THE TELESCOPE REQUIRES EXACTLY THE SAME SAFETY PRECAUTIONS AS FOR VISUAL OBSERVING – I.E., APERTURE FILTERS DESIGNED FOR SOLAR OBSERVING WITH TELESCOPES . Without the proper fi lters, the Sun is every bit as dangerous when seen in a camera’s view fi nder as it is when viewed with an eyepiece!

Your choice of aperture fi lter for photography is important, as some fi lters give much better results than others. Most fi lters for visual use have a “neutral density” of 5, which means that they transmit 1/100,000th, or 10 −5 , of the Sun’s light. Such fi lters are often known as ND5 fi lters. This usually gives an image that is bright enough for visual observing, even with quite high magni fi cations, but it can often be too dark for photography, particularly when you are shooting a highly magni fi ed solar image to capture details of sunspot groups.

A faint image needs a longer exposure to be recorded. But long exposures leave us prone to the problem of camera shake, which is discussed in the section on compact digital cameras. When the shutter in a DSLR camera is fi red, two things happen in quick succession: the shutter opens, and the mirror re fl ecting the image into the view fi nder fl icks upwards, out of the light path. Both actions, especially the mirror movement, cause vibrations. These are too small to be noticed when the camera is used for everyday photography with normal lenses, but at the magni fi cations used in telescopes even tiny vibrations are enough to blur the image.

Although some cameras have a smoother shutter release and re fl ex action than others (as will be explained below), all cameras produce some vibration when used with a telescope. The best way of getting around this problem is to use an exposure too short for any vibrations to register. Generally, exposures of 1/125 of a second or less come into this category. Such short exposures have the additional advantage that the effects of bad seeing conditions are greatly reduced. Obtaining a bright enough image for such short exposures to be possible with ND5 fi lters often used to be a problem when fi lm was the recording medium, because many of the best fi lms for solar photography were not sensitive enough. But today’s digital cameras easily have enough sensitivity to capture a whole-disc solar image with exposures of 1/500 of a second or shorter through an ND5 fi lter. If you fi nd that the Sun’s image on the LCD screen is too faint, you can try increasing the ISO setting on the camera; provided you keep to ISO 400 or below, increasing the ISO does not make the image signi fi cantly more “noisy.”

However, when you increase the magni fi cation to shoot close-ups of individual sunspots and groups, the image becomes fainter and the exposure required may be longer than 1/125 of a second. You could set your camera to ISO 800, 1600 or even beyond, but this signi fi cantly increases image noise, and in solar photography we want to capture as much fi ne detail as possible. The solution to the problem is to use a fi lter that transmits more light. The Baader AstroSolar fi lter, in addition to

147Filters

being a fi rst-rate visual fi lter, also gives a somewhat brighter image than other visual fi lters. It is excellent for whole-disc photography and can even be used at medium magni fi cations showing only part of the disc.

If you stick to low to medium magni fi cations, this type of visual fi lter may be all you need. But to obtain the best results at high magni fi cations you need a fi lter specially designed for solar photography. Baader make such a fi lter, metal-coated on polymer like their visual fi lters, called “AstroSolar Photo Film.” One version is claimed to have a neutral density of 3.8, meaning that it transmits 10 −3.8 , or just over 1/6,000 of the Sun’s light, nearly 17 times more than visual fi lters (Fig. 7.4 ). Newer versions of this fi lter claim a neutral density of 3.0, or a transmittance of 1/1,000 of the incoming light.

The “Type 3+” fi lter produced by Thousand Oaks Optical is an example of a glass photographic fi lter; it is coated to a neutral density of approximately 4 – i.e., it transmits 10 −4 of the Sun’s light, ten times more than visual fi lters. But these photographic fi lters are not safe for visual use, even for focusing the image through the camera view fi nder. Always focus the image using a visual fi lter fi rst, and then switch to the photographic fi lter when you are ready to make the exposure. With a photographic solar fi lter you can use exposures of 1/1,000 of a second or even shorter, even at higher magni fi cations. With a Baader ND3.8 fi lter on a 80 mm refractor, you can use an exposure of 1/1,250 of a second at ISO 400 on a Canon

Fig. 7.4 A fi lter made from Baader AstroSolar Photo Film, in a home-made mount and fi tted to the author’s 80 mm refractor. This type of fi lter transmits a brighter image than visual fi lters and so allows very short exposures to be used with DSLR cameras. All photographic fi lters should carry a warning label like this one, indicating that they are not safe for visual use

148 7 Imaging the Sun with a Digital Camera

300D to shoot close-ups of sunspot groups using eyepiece projection giving an effective focal length of nearly 5,800 mm and an effective f/ratio of f/72 (see below for details on photographic techniques). The same fi lter gives superlative results even at low magni fi cations for whole-disc shots – although then sometimes you might have to use an exposure of 1/4,000 of a second at ISO 100 – the shortest exposure and slowest speed available on this camera!

The prices of photographic fi lters are approximately the same per aperture as for visual ones. One small problem with the Baader ND3.8 material is that it is not available in A4-sized sheets, like its visual cousin, but comes in a roll 1 m long and 0.5 m wide. One solution to this could be to buy a roll collectively with other members of your local astronomy club.

Mounting the Camera

If your camera is of the compact type, it will not have a removable lens, and so how to position the camera to the telescope eyepiece requires some thought. You can get quite good results by simply holding the camera up to the eyepiece (Fig. 7.5a ), but you may well have problems with camera shake or partly obscured images caused by the camera being not quite squared on to the eyepiece. It is much better to mount the camera on a tripod (Fig. 7.5b ) – this both avoids vibration and allows you

Fig. 7.5 Three methods of using a compact digital camera at the telescope eyepiece to image the Sun. ( a ) A Canon PowerShot A640 camera held up to the eyepiece by hand. ( b ) The PowerShot A640 mounted on a tripod and pointed into the eyepiece. ( c ) A Nikon Coolpix 900 camera clamped to the eyepiece using a commercial digital camera adapter

149Mounting the Camera

to precisely center the camera over the eyepiece using the tripod adjustments. The tripod method is still inconvenient, though, as the telescope eyepiece and solar image are moving all the time relative to the camera, and so you frequently need to readjust the position of the camera on the tripod. This method is particularly awkward if you are shooting at high magni fi cation.

Best of all is to attach the camera to the telescope with some sort of bracket or adapter. A number of suppliers advertising in astronomy magazines have produced brackets that allow various models of cameras to be clamped to the telescope.

Fig. 7.5 (continued)

150 7 Imaging the Sun with a Digital Camera

Commercially made adapters come in two general forms. One type takes the form of a rod, at one end of which the camera is attached via its tripod bushing, and at the other is a clamp that attaches it to the telescope eyepiece or drawtube. A major advantage of this type of adapter is that it can be used with almost any digital camera, which is why it is sometimes known as a “universal” digital camera adapter. A disadvantage is that, like the hand-held or tripod methods described above, it lets in stray light between the camera lens and eyepiece.

Whether you use this type of adapter or the tripod or hand-held methods, it is a good idea to wrap a dark cloth around the area between the camera and the eyepiece, or even make a small tube of black paper to block out the stray light. The second type of adapter takes advantage of the fact that some cameras have lenses with threaded barrels. Some suppliers sell adapter rings for attaching these cameras to an eyepiece or camera adapter. Figure 7.5c shows a Nikon digital camera attached to the eyepiece using such a commercial adapter ring system. This avoids the problem of stray light entering the optical train between the camera and eyepiece, but the lenses on digital cameras come in many shapes and sizes and so you need to order a speci fi c ring for a speci fi c model of camera. Before buying a camera you may wish to enquire as to what kind of adapters are available for it.

If your camera is a DSLR, attaching the camera to the telescope is much easier, because this is done with standard camera adapters and rings that have been around since the days of 35 mm fi lm and are readily available from astronomical equip-ment suppliers. These adapters allow you to mount the camera very fi rmly and securely, and they allow no stray light to enter the light path between the camera lens and eyepiece. Which adapters to buy and how to use them are discussed in the sections below on prime focus and eyepiece projection photography.

Photographic Techniques

Shooting the Projected Image

Solar photography can be almost as simple as everyday photography. If you use the projection method to observe the Sun, you can photograph the projected image quite easily. As the projected image is brighter than its surroundings, you will get the best results if you use a slightly shorter exposure or slightly smaller f/stop than you would for normal photography in sunlit conditions.

Shooting the projected image has two major problems, however. The fi rst is that you can never photograph the Sun’s image from directly above, since then the eye-piece would be in the way. Therefore the solar disk will be somewhat distorted in the photograph, due to the effect of perspective. Secondly, it is dif fi cult to get good contrast in photographs of a projected image, and sunspot detail often appears washed out. The resulting photograph also tends to be rather “messy,” with tele-scope parts and other objects intruding into the picture. You can correct the distor-tion by placing the camera as close as possible to the telescope drawtube and tilting

151Photographic Techniques

the projection screen towards the camera. Tilting, however, will throw parts of the solar image slightly out of focus. An alternative could be to scan the image and use digital processing techniques to eliminate the distortion. You can also increase the contrast by projecting the image onto a blank wall or perhaps a slide projection screen in a darkened room. However, almost no photograph of a projected image even comes close in quality to an image shot directly through a telescope. Photographing pro-jected images at high magni fi cation is especially dif fi cult, as the contrast is then even lower. The technique is useful in emergencies, though – for example, when a projected image showing a large sunspot is on display at an astronomy club gathering and you are unable to photograph the Sun with your own telescope.

The Afocal Method

The simplest way of taking pictures of the Sun is to point the camera, with its lens in place, into the telescope eyepiece. This is called the afocal method, and if your digital camera is of the compact type this is the only way you can photograph the Sun directly through a telescope, since you cannot remove the lens on this type of camera. The afocal method has long been popular with amateur astronomers as a way of photographing the Moon and the brighter planets with undriven telescopes that are dif fi cult to attach cameras to, such as re fl ectors on Dobsonian mounts. As described in the section on mounting the camera, it is best to use a tripod or an adapter to mount the camera at the eyepiece, so as to ensure your images are not blurred by camera movement.

The afocal method has some disadvantages. Although it is quite easy to center and focus the image when you use a low- to medium-power eyepiece, both become dif fi cult at higher magni fi cations, especially as the image on the camera’s LCD view fi nder can be dif fi cult to see in the sunlight. Another problem, already noted above, is stray light between the camera and the eyepiece. You can eliminate this with a piece of dark cloth or a tube of black paper between the camera lens and eyepiece, but unless you use an adapter, the telescope is moving all the time relative to the camera and so will eventually drift out of position!

The afocal method can be made to work very well with a compact camera if you mount the camera correctly and use other techniques for taking good pictures, described under “Taking Pictures.” If your camera is a DSLR, it may be worth try-ing this method as a fi rst experiment, but because you can remove the lens on a DSLR, you have a choice of much better methods with this type of camera.

Prime Focus Photography

In the prime focus method, the eyepiece is removed from the telescope, the lens is removed from the camera and the camera body is attached directly to the telescope

152 7 Imaging the Sun with a Digital Camera

(Fig. 7.6 ). Effectively, this method uses the telescope as a long telephoto lens. It is, in fact, possible to get reasonable images of the Sun using a very long telephoto lens system intended for terrestrial photography, say a 500-mm telephoto with a 2× teleconverter – always provided that the front of the lens is covered with a safe aperture fi lter. Such a setup can provide shots showing the Sun’s full disc with the larger sunspots. But to take detailed sunspot photographs you need a telescope.

When photographing the Sun (or the Moon, which has the same angular size as the Sun) astrophotographers use the following formula to determine the diameter of the Sun’s disc, in millimeters, on the image sensor:

110

f

where f is the focal length of the telescope. So in a telescope of 1,000 mm focal length, the size of the prime focus image would be:

1,000

110

9.1 mm=

Fig. 7.6 Prime focus photography with a DSLR: the camera body is attached directly to the telescope drawtube with an adapter

153Photographic Techniques

Consumer DSLRs today tend to have image sensors measuring 15 mm × 22 mm, so a 9.1 mm solar image will fi ll a good part of the picture and show all the major sunspots. Indeed, if you are shooting a total solar eclipse with a DSLR, 1,000 mm is about all the focal length you will ever need, since you will want to include the corona and some background sky in the picture as well as the Moon’s silhouette covering the Sun. Shooting at prime focus is also an excellent way to record a partial solar eclipse with a small telescope (Fig. 7.7 ). Only if your camera is an expensive “full frame” model – i.e., one with an image sensor the same size as 35 mm fi lm (24 mm × 36 mm) – will you want a somewhat longer focal length.

Prime focus photography gives excellent whole-disc shots showing the appear-ance of the Sun on a particular day. For example, with a Meade ETX 90 mm Maksutov telescope (focal length 1,250 mm), shooting at prime focus gives a solar disc about 11.4 mm across, capable of showing plenty of detail. A “classic” 150 mm (6 in.) f/8 Newtonian re fl ector gives a 10.9 mm image. But another “classic” amateur’s telescope, a 200 mm (8 in.) f/10 SCT, of 2,000 mm focal length, gives an image 18.2 mm across. This is larger than the image sensor’s smallest dimension in many DSLRs, and the resulting picture will show two limbs of the Sun “lopped off”. Fortunately, the largest manufacturers of these instruments, Meade and Celestron, offer devices to reduce the telescope’s focal length and so allow these instruments to photograph the whole disc at the prime focus.

Fig. 7.7 Prime focus photography can be a good method for imaging a partial solar eclipse. This image of the partial eclipse seen just after sunrise from the UK on January 4, 2011, was taken by the author using a Canon 300D at the prime focus of a Vixen 80 mm refractor (focal length 910 mm), with a Baader AstroSolar Photo Film (ND3.8) fi lter. Exposure 1/125 s at ISO 400

154 7 Imaging the Sun with a Digital Camera

To take prime focus shots, you need to couple the camera to the telescope using a camera adapter designed for this purpose. The adapter itself attaches to the camera using a T-ring , a device with a bayonet fi tting to slot into the camera like a lens. The front of the T-ring has a standard thread that should be the same as the thread on the adapter. As with camera lenses, T-rings vary according to the make of the cam-era, and you need to get the one that fi ts your camera. T-rings used to be easily obtainable from good camera stores, but they are not as easy to fi nd nowadays, as their main purpose used to be for stacking lenses on top of each other to increase the focal length, and so they have become redundant in an age when good quality zoom lenses are readily available. Fortunately, you can buy T-rings for a number of different camera makes from many of the astronomical equipment suppliers adver-tising in astronomy magazines.

The same suppliers also sell camera adapters. The type of adapter you need, however, depends on the type of telescope you are using. If you have a refractor or a Newtonian, you need a standard camera adapter, the front of which pushes into the drawtube like an eyepiece. This type of adapter unscrews into two sections: a T-shaped front section and a long cylinder at the rear. You need just the front section for prime focus work. This screws into the T-ring, which in turn attaches to the camera, and the whole system is then inserted into the drawtube. If your telescope is an SCT or Maksutov you need a different type of adapter, which is known, rather confusingly, as a T-adapter . The rear threads into a camera T-ring, as before, but the front threads onto the rear cell of the telescope. Meade and Celestron T-adapters work with any of the SCT range produced by these companies, but if your telescope is one of the Meade ETX range you need a special T-adapter made for this line of telescopes.

The prime focus method is fi ne for whole-disc photography, but with most instruments it does not provide enough magni fi cation to show fi ne details within the spots and in no instrument does it exploit the full resolving power of the telescope. Also, with some telescopes, such as small refractors and short-focus Newtonians, the image formed at the prime focus is rather small (Fig. 7.8 ). An 80 mm refractor, for example, has a focal length of 910 mm, giving a prime focus image just 8.3 mm in diameter, little more than half of the shortest dimension of a DSLR image sensor.

To magnify the image further there are three somewhat more advanced methods for photographing the Sun.

Using a Teleconverter

You can magnify a prime focus solar image by a modest amount quite easily with an ordinary teleconverter lens between the camera body and the telescope. A tele-converter is especially useful if your telescope’s focal length is a little too short to give a good-sized prime focus image. For example, a 1.4× teleconverter effectively increases the focal length of an 80-mm refractor (focal length 910 mm) to 1,274 mm, which produces a disc size of 11.6 mm, fi lling much of the frame in a DSLR.

155Photographic Techniques

Figure 7.9 is an example of one of the author’s own photographs taken using a teleconverter. If your telescope has a very short focal length – say 600 mm or less – you may want to use a 2× teleconverter to effectively double the focal length.

Note, however, that as well as increasing the effective focal length, teleconverters increase the f/ratio as well. Because f/ratios follow a square law, doubling the f/ratio requires an exposure four times as long. For example, if the Sun is correctly exposed at 1/2,000 of a second at prime focus, you will need a 1/500 s exposure if you use a 2× teleconverter. You can stack two or more teleconverters to give even longer focal lengths, although this increases exposure times even further.

The main problem with using a teleconverter with a DSLR on a telescope is that modern camera lenses and teleconverters are designed to operate electronically as a system with the main camera, so that focus, aperture and shutter speed can be determined automatically. This means that when a camera is used with a telecon-verter in front of it, but no lens, it thinks there is something wrong with the system. When you use a 1.4× teleconverter with a Canon 300D on a telescope, you get an error message when you press the cable release and the shutter refuses to fi re! You can get around this by attaching the teleconverter so that it does not click into place. This way it is mounted on the camera without being electronically connected to it. Take great care when doing this, however, as the camera body is not securely con-nected to the teleconverter and telescope and so could possibly fall to the ground. Do not use this method if your telescope is set up on a hard surface!

Fig. 7.8 A small telescope of short focal length gives only a very small solar image when the camera is used at the prime focus. Image taken by the author on March 14, 2010, using a Canon 300D DSLR at the prime focus of a Takahashi FS-60C 60 mm refractor, focal length 355 mm, equipped with a full-aperture Baader AstroSolar (ND5) fi lter. Exposure 1/2,000 s

156 7 Imaging the Sun with a Digital Camera

An alternative to teleconverters is a Barlow lens, since a teleconverter is a negative lens, like a Barlow. However, some cheap Barlows can cause internal re fl ections. Some otherwise good solar images taken through a low-cost achromatic Barlow can be spoiled by such a re fl ection.

Eyepiece Projection

To take close-ups of individual sunspots and groups you need a high magni fi cation and therefore a long effective focal length. The best way to achieve this is to use eyepiece projection . Again, the camera lens is removed and an ordinary telescope eyepiece projects an image of the Sun onto the fi lm – not unlike projecting the image onto a screen for visual observing, although for photography we must, of course, use an aperture fi lter. If your telescope is a refractor or a Newtonian this is where the second, cylindrical part of the adapter comes in. An eyepiece is inserted into a receptacle inside this section and locked in place with a thumbscrew. This section is screwed to the front part of the adapter and the whole adapter is then threaded into the T-ring (see the diagram in Fig. 7.10 and the photograph in Fig. 7.11 ).

Fig. 7.9 Whole-disc solar image, showing a large sunspot group. Image taken by the author on September 13, 2005, using a Canon 300D DSLR on an 80 mm f/11.4 refractor with a 1.4× teleconverter, giving an effective focal length of 1,274 mm. The Sun’s disc nicely fi lls the frame in many DSLRs using this sort of focal length. Filters used were Baader full-aperture AstroSolar Photo Film (ND3.8) and No. 8 (light yellow) secondary fi lter. Exposure was 1/3,200 of a second at ISO 100

157Photographic Techniques

Telescope Drawtube

Camera Adaptor

Sun’s image projectedonto imaging sensor

DSLR Camera

Eye pieceT-Ring

Fig. 7.10 Diagram showing the principle of eyepiece projection using a refractor or Newtonian telescope, using a camera adapter and T-ring

Fig. 7.11 Eyepiece projection arrangement on the author’s 80 mm refractor, showing Canon EOS 300D DSLR (with cable release), T-ring and camera adapter

158 7 Imaging the Sun with a Digital Camera

If your telescope is an SCT you need a tele-extender , a simple tube that threads into the T-ring at the rear and into the telescope’s rear cell at the other end. The amount of magni fi cation produced by eyepiece projection depends on the focal length of the eyepiece used and the distance between the eyepiece and the sensor. Higher magni fi cations are given by shorter focal length eyepieces, as in visual observing, or by increasing the eyepiece-to-sensor distance. A simple camera adapter allows you to change the eyepiece-to-sensor distance by just a few millimeters at best, and then only by sliding the eyepiece in its receptacle. It is better to use a slightly more expen-sive “variable projection” adapter, which uses two tubes sliding inside each other to vary the projection distance by 30 mm or more.

To determine the effective f/ratio given by eyepiece projection, we need to take into account the eyepiece focal length and the projection distance in addition to the tele-scope’s f/ratio. We can use the following formula to work out the effective f/ratio:

( )telescope eyepiece

eyepiece

f d fl

fl

´ -

where telescopef is the telescope’s original f/ratio, d is the projection distance and

eyepiecefl is the focal length of the eyepiece. As an example, let us suppose we want to get a close-up of a sunspot using an

80 mm f/11.4 refractor and a 15 mm eyepiece. Variable projection adapters gener-ally allow a range of projection distances ( d ) between 110 and 150 mm. Let us select a mid-way value of d , 130 mm. The projection distance is measured between the focal plane of the eyepiece and the plane of the image sensor. In practice, it is impossible to measure this exactly, because the focal plane varies from eyepiece to eyepiece and you cannot see the eyepiece inside the adapter. Estimate the distance by measuring from where you judge the mid-point of the eyepiece to be. The posi-tion of the image sensor plane should be marked by a small circle with a line through it, engraved in the top or bottom of the camera. Plugging the numbers into the formula, we get:

( )11.4 130 15

15

´ -

/ 87.4f=

We can fi nd the effective focal length provided by this setup by multiplying the new f/ratio by the telescope’s aperture. Multiplying 87.4 × 80 gives us 6,992 mm – a long focal length, showing only a small part of the Sun’s disc, excellent for detailed images of individual sunspot groups. Figure 7.12 shows an image of two sunspot groups taken through the author’s 80 mm refractor using eyepiece projection.

If your telescope has a very short focal length, you can in theory use eyepiece projection to obtain whole-disc shots of the Sun. In practice, however, eyepiece projection is not ideal for this purpose. To get the Sun’s full disc in the picture you may have to use a low-power eyepiece and a short projection distance, which can

159Taking Pictures

cause the fi eld to be “curved” – that is, the center of the picture is in focus but the periphery is blurred. If the curvature is small you may be able to get away with it, though, as the limb darkening on the Sun can cause the edge of the disc to be under-exposed, and a small amount of blurring may not be noticeable! But generally the best way of increasing the image size given by a short-focus telescope is to use a teleconverter or Barlow lens as described above.

The results you can achieve with eyepiece projection are also limited in that with the seeing commonly encountered in the daytime you will seldom obtain images that show detail down to the resolution limit of your telescope. Moreover, your chances of obtaining such images decreases with the aperture of the telescope, so large telescopes almost never give images down to the resolution limit with com-pact or DSLR cameras. You can increase your chances of getting really sharp images by taking lots of pictures and selecting the best one (see “Taking Pictures” below), but the best way of obtaining the highest-resolution images is by using a webcam or other movie-type digital camera. How to do this is described in Chap. 8 .

Taking Pictures

A golden rule in solar photography (and any other astronomical photography) is always to take notes . Only with careful notes made at the time you took the pictures can you determine what is the best combination of exposure, f/ratio, fi lter and other

Fig. 7.12 Close-up of two sunspot groups, an enlargement of an image photographed by the author on December 4, 2005, using an 80 mm refractor and eyepiece projection with a 15 mm eyepiece to give an effective focal length of 5,763 mm at f/72. Filters used were Baader full-aperture AstroSolar Photo Film (ND3.8) and No. 8 (light yellow) secondary fi lter. Exposure was 1/1,250 s at ISO 400 on a Canon 300D DSLR

160 7 Imaging the Sun with a Digital Camera

factors. It is essential to give full details of how you took each image if you send pictures to solar observing organizations or to magazines for possible publication. In addition to the exposure, ISO speed, focal length, f/ratio and fi lter, for each exposure you should also record the date and time (in UT) at which you took it, together with the observing conditions. It is also important to note what method you used (e.g., eyepiece projection or prime focus), any means of reducing vibration and any problems encountered. The exposure data that your digital camera records with each image is not suf fi cient, because it does not record, for example, what kind of seeing conditions you took the image in. The camera also omits what focal length and f/ratio you used when your camera is attached to the telescope; when you remove the lens, a DSLR normally records the f/ratio as “00”. In any case, most these data are lost when you delete an image from the memory card.

Seeing conditions are very important in solar photography. Poor seeing can cause even whole-disc shots, taken at relatively low magni fi cations, to be blurred and can render close-ups of sunspots useless. The advice given in Chap. 3 on seeing for visual observing applies to photography as well. If you use the projection method for visual work, you may wish to do any high-resolution photography beforehand, as projection allows heat to build up inside the tube, causing turbulence if you do photography just afterwards. Thin, high clouds such as cirrus can also dramatically reduce the contrast of solar detail on photographs, even though you can often carry on visual observations in such conditions. But even if the seeing is poor or signi fi cant cirrus is present, it is remarkable what can be achieved with digi-tal imaging techniques. The very short exposures used for the Sun with today’s digital cameras can capture precious moments of better seeing and it is also possi-ble to increase the contrast of images using the image processing techniques described in Chap. 8 .

Before starting a photographic session, make sure that any eyepieces or other lenses you use, such as teleconverters, are scrupulously clean and free from dust specks. These lenses are often located close to the focus of the telescope, and dust particles show up on images as gray blobs. Dust blobs are not usually noticeable on night-sky pictures, but they stand out very distractingly against the bright background of the Sun and can sometimes be confused with sunspot penumbrae!

The majority of dust forms on the exterior surfaces of lenses, and you can clean them quite easily using a lens cleaning kit from a camera store. Always use a blower brush (usually provided with the kit) fi rst to remove any large dust particles, as these can scratch the surface of the glass if you try to wipe the lens straightaway. Over the years a small amount of dust seeps into the inside surfaces of the lenses. If a signi fi cant amount accumulates you may have to take the lens to an astronomi-cal supplier or a camera repair specialist to have it cleaned professionally. Do not take eyepieces apart unless you know what you are doing, as it is easy to insert lens elements back in the wrong way around.

If your camera is a DSLR, a much more serious problem is dust on the image sensor itself – or rather, on the infrared blocking fi lter just in front of the sensor.

161Taking Pictures

These show up as distracting dark spots on the image. Dust particles are far more of a nuisance to the solar photographer than they are in terrestrial photography, for two reasons. The main reason is that terrestrial photography typically uses lenses at fast f/ratios – normally f/8 and sometimes even faster – and so any dust particles are blurred to the point where they are not noticeable. But in solar photography we are often using much longer f/ratios: f/20 or f/60 are typical and even f/100 is not uncommon if we are imaging the Sun at very high magni fi cation, and at these high f/ratios the dust particles are sharp, with even small ones being noticeable. Secondly, in solar photography we are imaging against a smooth, bright back-ground, against which dust particles show up much more starkly than they do against everyday terrestrial scenes and even more so than against the night sky. The problem is made even worse by the fact that dust particles appear as small, slightly ill-de fi ned dark spots – just like small sunspots!

Unfortunately, there is no easy solution to the problem of dust on the imaging sensor. DSLRs have a sensor cleaning mode, which raises the camera’s internal mirror and holds open the shutter, giving access to the sensor for cleaning. But you must take great care when cleaning the sensor, as touching or wiping it can scratch it or cause any dust particles present to make scratches. For this reason, you should not use a “blower brush” designed for cleaning camera lenses, as this could cause dust particles to scratch the sensor and may itself introduce more dust to the sensor. A better idea is to use a “hurricane blower,” an air bulb that blows compressed air through a nozzle, blowing the dust particles off the sensor without actually touch-ing it. But even this type of device might do no more than blow dust around the interior of the camera. Always use a blower with the imaging sensor facing down-wards, as then the force of gravity will cause any particles blown off to fall towards the ground and not back on to the sensor. Some modern DSLRs have a built-in dust reduction feature that works by vibrating the sensor ultrasonically.

The best solution to the dust problem is to do your best to prevent dust from getting onto the sensor in the fi rst place. To this end, keep the lens or camera body cap on the camera at all times when you are not using the camera, and store the camera in a high-quality case in as dust-free an environment as possible. When changing lenses, keep the camera body pointing downwards as much as possible, to prevent dust from falling onto the sensor. Obviously this is not possible when inserting the camera into a telescope aimed at the Sun, but you can still assemble the various adapter parts required with the camera pointed down before it is placed in the drawtube. It is also a good idea to thread a transparent fi lter of some kind into the front of the adapter, to prevent dust entering the camera while it is in the tele-scope – indeed, as described later in this section, such secondary fi ltering will improve your photographs anyway.

As is explained in Chap. 8 , it is possible to remove some dust spots using a digi-tal image processing technique – though you need to ensure that what you are removing really is dust spots and not small sunspots! Also, if you send your images to a solar observing organization, or have them published in print or online, you should always state that the images have had dust removed at the processing stage, as any such retouching of an image can compromise its scienti fi c accuracy.

162 7 Imaging the Sun with a Digital Camera

Taking Pictures with a Compact Digital Camera

If you are using a compact digital camera, begin by carefully focusing the fi ltered solar image in the eyepiece. If you wear glasses, it is important to focus the telescope with them on, as without them the focus will be inaccurate. Then mount the camera on the tripod or adapter you have chosen. Switch on the LCD and ensure that the camera is aimed squarely into the eyepiece. If you are using a tripod, and if your telescope is a refractor or an SCT, turn the star diagonal around so that the eyepiece is in a horizontal position. This makes positioning the camera much easier than when the diagonal is in its usual upright position. Also, place the camera a centimeter or so away from the eyepiece to allow the drawtube to be racked out for focusing. The camera’s distance from the eyepiece does not affect the focus in afocal photography, so you can move the camera closer in once you have fi nished focusing.

Use a low-power eyepiece to begin with. Higher powers have less “eye relief” – that is, you need to place your eye or camera lens closer to it to see the full fi eld of view, which makes it dif fi cult to align the camera. Choose a magni fi cation that shows the whole solar disc comfortably within the fi eld of view. If your camera has no manual exposure facility, avoid using a very low power that shows the Sun’s disc surrounded by a lot of black sky. A camera relying on automatic exposure determines the shutter speed from the average brightness of the picture. Too much black sky can give a low average brightness and so cause the camera to overexpose the Sun’s disc. Many auto-exposure cameras do have exposure compensation and spot metering facilities, but to begin with it is simpler to use automatic expo-sure. One way of increasing the magni fi cation enough for the Sun to fi ll most of the frame is to use a low-power eyepiece with a Barlow lens, which boosts the power but does not sacri fi ce eye relief. Using a Barlow lens with a medium-power eyepiece is also a good way to achieve high magni fi cations when you are ready to go on to more advanced work, as you still have some eye relief for aligning the camera. Best of all, use the zoom facility on your camera’s lens. Remember not to use “digital zoom” – as noted above, this merely increases the size of the pixels and does not genuinely increase the magni fi cation.

The LCD screen can be hard to see in full sunlight, even when turned up to its maximum brightness. This is especially true if you are using a Newtonian or a star diagonal and the Sun’s light is striking the screen sideways-on. To make it easier to see, try tilting the screen away from the sunlight, if it can be tilted with respect to the camera, or you could make a simple shade from black card and stick it to the camera with tape. It is also important to block out, or at least minimize, stray light between the camera lens and eyepiece. As noted above, an adapter ring attaching the camera to the eyepiece neatly solves this problem. Otherwise, wrap a dark cloth loosely around the camera and drawtube, ensuring it does not prevent you from using the camera’s controls, or make a short tube from black paper so that it hangs loosely over the eyepiece and lens barrel and bridges the gap between them. Make sure it is not too tight, to avoid transmitting camera vibrations to the telescope and to allow for minor adjustments of the camera’s position.

163Taking Pictures

You now need to ensure that the image on the LCD screen is focused. Most digi-tal cameras have automatic focusing – commonly referred to as autofocus – and, as with exposure, many popular models have no means of overriding this and setting the focus manually. Fortunately, many digital cameras focus quite accurately on the solar image in the eyepiece, as long as you have focused it carefully with your eye beforehand. Many observers fi nd it hard to determine whether the image on the LCD is precisely in focus, not only because stray sunlight makes the screen look dim but also due to the small size of the image and the screen’s rather low resolu-tion. If your camera has a video output socket, you could connect it to an external monitor or a portable TV and so view the image on a larger and brighter screen. A number of amateurs have used this method successfully, but if your telescope setup is portable it is inconvenient to have to set up a TV for each imaging session in addition to the telescope and camera. If you have any doubt as to the focus of the picture, it may be better to use trial and error in your fi rst imaging sessions, by marking the telescope drawtube with several focus positions and taking a series of pictures with the drawtube in a different position each time. If you make a note of the focus position for each picture, you can establish which position gives the best focus when you view the images on the computer later.

Even if you have to rely on auto-exposure, you can control the brightness of the image to some extent using the ISO rating. You may fi nd that a slow setting, such as 100 ISO, is quite adequate for solar photography, as the Sun’s disc is a bright subject, even when imaged through an aperture fi lter. Remember also that a higher ISO gives a more “noisy” picture. In fact, image noise at high ISO speeds is more of an issue with compact cameras than it is with DSLRs. Only if you are shooting at high magni fi cation or in H-alpha or Ca-K should you select a faster speed to keep exposure times down.

A fi nal important setting is the quality of the image. Remember that this refers to the depth of the image (i.e., the extent to which the data is compressed) and not the resolution in megapixels. The higher the quality of the image, the greater its depth and so the more realistic are the tones. Compact cameras generally allow you to choose a “small,” “medium” or “high” image quality, and many of them offer some median options in between these three. Some – though not all – compacts allow you to select an uncompressed format, such as TIFF or the camera’s propri-etary RAW format, as the highest image quality. Although the highest quality gives the best picture, it is better to select a “high” quality JPEG setting to begin with or if you are experimenting with new techniques. Uncompressed images take up a lot of memory and may quickly fi ll up your card and take a long time to download to your computer. Taking a large number of high-quality JPEG images allows you to experiment widely without changing cards and to download the results to the com-puter more quickly. You can take a few RAW images when you are satis fi ed that you have got the alignment, focus and exposure correct. In any case, you will be pleasantly surprised at how much sunspot detail you can record even on high-quality JPEG images.

When you have aligned the camera, focused the image and adjusted the various settings, you are ready to expose. Squeeze the shutter release button as gently as

164 7 Imaging the Sun with a Digital Camera

you can. Some digital cameras require you to press the shutter release quite hard, and the camera can move by a surprising amount, even when it is mounted on a separate tripod. Always use the self-timer facility or remote control if your camera has either of these features. Some amateurs have made their own ingenious brack-ets on which to mount a cable release, and it is even possible to buy such cable release brackets commercially from one or two sources.

Be prepared for a high failure rate to begin with, as there are so many things you have to get right in digital photography with the afocal method – controlling stray light, aligning the camera, magni fi cation, focusing, exposure and camera shake. However, the main dif fi culty in afocal digital photography usually occurs if the camera is mounted separately from the telescope, as it is then dif fi cult to maintain camera alignment during the exposure.

Because the available fi eld of view is often restricted, even a small movement of the camera, such as that caused by pressing the shutter release, can be enough to jog the camera so that the Sun is half-hidden. Even the self-timer does not always compensate for this problem. This is why a camera capable of being attached to the eyepiece or, better still, a digital SLR, is much preferred for solar photography. Some cameras allow you to magnify the image on the LCD screen once you have taken it and examine parts of it in close-up. If your camera has this feature make use of it, as the low resolution of the LCD monitor can hide minor fl aws in an image. But whatever camera you have, digital photography’s instant results and the ability to throw away an image there and then if you don’t like it, allows you to experiment to your heart’s content. You can take 10 or 20 “dud” images and then have the satisfaction of getting a good one. Don’t be too hasty about throwing away images, though, unless they are really bad, as if it looks reasonably well-centered and focused you can usually do something with it at the image processing stage. As will be explained in Chap. 8 , it is surprising what you can do even with a relatively poor image.

Taking Pictures with a DSLR

The fi rst step in taking pictures of the Sun with a DSLR is to attach the camera to the telescope using the appropriate combination of adapter and T-ring as discussed earlier under “Photographic Techniques.”

Focusing the Sun’s image with a DSLR is much easier than with a compact camera, since you can focus on the image in the camera’s optical view fi nder, much as you would when observing visually through a telescope eyepiece, rather than using the somewhat fuzzy, low-contrast image on a compact camera’s LCD screen. But achieving a precise focus, which makes full use of your camera’s and tele-scope’s resolution, requires great care. The through-the-lens focusing systems on DSLR cameras work (as did their forebears on 35 mm SLRs) by projecting an image onto a special focusing screen in front of the pentaprism that re fl ects the image into the view fi nder. These focusing screens are designed for everyday

165Taking Pictures

photography using ordinary camera lenses at relatively fast f/ratios, not astronomical photography with very long f/ratios. As a result, the screens have a coarse texture and the telescopic image of the fi ltered Sun on a DSLR focusing screen appears dim, and it can be dif fi cult to fi nd the exact focus.

The problem gets worse as you increase the magni fi cation. In the age of 35 mm fi lm, some of the best fi lm DSLRs solved the problem by allowing the photographer to substitute the standard focusing screen with a screen containing a clear central section on which was engraved a fi ne cross, allowing your eye to focus at the correct point. Some manufacturers of astronomical accessories even sold special bright screens for astrophotography. But only the most expensive professional DSLRs have interchangeable screens, leaving most of us stuck with the standard screen that comes with the camera. Some modern DSLRs have a “live view” mode that allows you to focus using the LCD display at the rear of the camera. This can be advantageous in night-sky photography through a telescope, when the image is too faint to see at all through the optical view fi nder, but for the solar photographer it offers none of the advantages of through-the-lens focusing with a DSLR and all the disadvantages of focusing using the LCD screen with a compact camera. A third alternative is to focus the image remotely using a computer, or a laptop if out in the fi eld, as many DSLRs come with software that supports remote control of the camera by a computer. If your telescope is permanently set up, this may be a good way to go, but if your telescope is portable you have the inconvenience of setting up a computer with the camera every imaging session.

A preferred method of focusing a DSLR on the telescopic solar image uses the problem of the rough-textured focusing screen to solve itself. First of all, it is essen-tial to remember that the image we see through the camera’s optical view fi nder is not the visual image that we see through a telescope eyepiece but an image pro-jected onto a screen, just like the image projected onto the camera’s image sensor. Therefore, we can only be sure that the image is in focus when the screen itself is in focus. The screen can be focused independently using the camera’s diopter adjustment, which takes the form of a very small wheel by the side of the optical view fi nder. (You may need to remove the rubber eye guard over the view fi nder to gain access to it.) To focus the screen accurately, the screen needs a reasonably bright and, above all, even background illumination. This can be achieved by grossly defocusing the fi ltered solar image through the telescope so that you get an amorphous white background. For even better results, you should stop down the aperture of your telescope, as this gives a narrower incoming optical beam, which produces sharper shadows, so that when the coarse focusing screen is in focus its structure shows up sharply as a fi ne pattern.

If your telescope is a refractor, you can often stop down the aperture quite easily by removing the secondary cap in the center of the objective lens cap, a feature found on many modern refractors. When you see a sharp image of the focusing screen, you can be con fi dent that the solar image, however faint, is correctly focused when it appears sharp. When focusing the solar image, avoid focusing on the outer part of the screen, partly because the image is faint here but mostly because the center of the image will be blurred if there is any fi eld curvature.

166 7 Imaging the Sun with a Digital Camera

For both focusing the screen and the solar image, it is very helpful to use a right-angled focusing magni fi er such as the “Angle Finder C” supplied by Canon and designed for use with their line of DSLRs (Fig. 7.13 ). Such a device is not cheap, but it is one of the most useful accessories you can buy for solar photography with a DSLR, as the extra magni fi cation it provides allows you to focus much more accurately. Moreover, it allows you to view the Sun’s image at a much more com-fortable angle, which in itself helps you to focus more accurately. In fact, it makes photography more enjoyable because you can look at the Sun’s image much as you would with an eyepiece or star diagonal. The image of the partial eclipse of March 29, 2006, shown in Fig. 7.14 was taken with the help of this focusing magni fi er and the focusing method described here.

Before starting exposures, carefully center the image of the Sun or your target sunspot group in the camera’s view fi nder. If your telescope has no motor drive, posi-tion the Sun’s disc or the sunspot you are photographing a little to the east of the center of the screen, to allow for movement of the image due to Earth’s rotation in the seconds before the exposure. Use trial and error to determine how far you have to position your target. To achieve consistent results, you should also orientate the camera correctly with respect to north. You can do this quite easily using the markers on the screen showing the position of the camera’s autofocus points. Two or three of these should be aligned horizontally. Adjust the telescope’s slow motion controls so that a small sunspot is placed at the top or bottom edge of the easternmost marker.

Fig. 7.13 The Canon “Angle Finder C” focusing magni fi er fi tted to a Canon 300D camera mounted on the author’s 80 mm refractor

167Taking Pictures

Switch off the telescope’s motor drive (if your telescope has one) and let the sunspot drift. Carefully adjust the orientation of the camera in the drawtube until the sunspot drifts onto the next marker along. When the sunspot’s direction of drift is aligned accurately to the markers, you know that the image is correctly oriented east–west. This technique is the photographic equivalent of orienting the projection screen for visual observing. This is important when sending in your images for analysis and, if done accurately enough, allows your images to be measured for the purpose of deriving sunspot latitudes and longitudes.

Trial and error is necessary for determining exposure as well. Filters, f/ratios and conditions vary so much that it is impossible to suggest precise exposure times. To begin with, take a large range of exposures starting with your camera’s fastest shutter speed (often 1/4,000 of a second) and working upwards, making careful notes and reviewing each image afterwards on the LCD screen.

Even when you have found the correct exposure, always take an exposure one stop shorter and another a stop longer than the ideal. This is known as “bracketing” exposures and allows for variations in the Sun’s brightness due to observing condi-tions or the altitude of the Sun. In fact, it is a good idea to take two shots at each shutter speed, so that you cannot confuse dust particles, scratches or faulty pixels with real solar features. Exposures of 1/1,000–1/2,000 of a second are excellent for whole-disc shots with a small refractor fi tted with a full-aperture Baader AstroSolar

Fig. 7.14 Partial solar eclipse, March 29, 2006, focused using the method for focusing a DSLR described in the text. Image taken with Canon 300D on 80 mm refractor with full-aperture Baader AstroSolar Photo Film (ND3.8) fi lter and 1.4× teleconverter. Exposure 1/3,200 of a second at ISO 100

168 7 Imaging the Sun with a Digital Camera

visual fi lter (neutral density 5). In practice, though, you might want to use a photographic fi lter (Baader AstroSolar Photo Film, ND3.8) even for whole-disc work with an a small telescope. When the Sun is at a reasonably high altitude, 1/4,000 or 1/3,200 s at ISO 100 gives excellent whole-disc shots with this fi lter. For eyepiece projection close-ups of sunspots with the same telescope and fi lter, try a 1/1,250 of a second at ISO 400.

If you use a Mylar-type fi lter, you may fi nd that the contrast of your images is improved if you use a yellow fi lter in addition to the main aperture fi lter. The use of such “secondary fi ltering” for visual observing was mentioned in Chap. 2 , and the same light yellow (number 8) fi lter as that used visually with the Baader AstroSolar aperture fi lter also works very well for photography, especially whole-disk images. This not only gives DSLR images a more realistic color, but it also increases the contrast of sunspots and slightly reduces the effect of chromatic aber-ration. Even better results can be obtained with the “Fringe Killer” fi lter produced by Baader Planetarium. This is an interference fi lter that works on the same prin-ciple as the expensive narrowband fi lters for observing the chromosphere, albeit with a much broader passband: it simply cuts off the red and blue/violet extremes of the spectrum and lets the rest of the light through. It noticeably reduces chro-matic aberration on photographs, especially images taken at high magni fi cation, and it also increases contrast, especially on granulation and faculae. Some ama-teurs, especially those imaging the Sun using Herschel wedges, have also had suc-cess with Baader’s “Solar Continuum” fi lter. This is also an interference fi lter, but it does not merely reduce chromatic aberration: its passband is optimized speci fi cally for showing photospheric detail. All these fi lters are designed to screw into a standard 1¼-in. eyepiece. Camera adapters have the same internal threads as eyepieces and you can screw your light yellow or Baader Fringe Killer fi lter into the front of the adapter. These fi lters let through most of the incoming light and so exposure times are hardly affected. All three of these fi lters are relatively inexpen-sive and are well worth experimenting with.

Photographing the Chromosphere

Photographing the chromosphere and its features in H-alpha (and Ca-K) is more dif fi cult than white-light photography, partly because much less light is available and also because the contrast of chromospheric features is much lower than sun-spots and other features of the white-light Sun. With off-the-shelf compact digital cameras and DSLRs, it is dif fi cult to take close-ups of prominences and other fea-tures at very long effective focal lengths as you can in white light, unless you have a large telescope, because the image becomes all but impossible to focus accurately and the exposure times are very long, resulting in blurred images.

The main methods for photographing the chromosphere are the afocal method, if you have a compact camera, and the prime focus method if you use a DSLR. If your telescope has a very short focal length (as is often the case with small

169Photographing the Chromosphere

apochromatic refractors), you can use a teleconverter or Barlow lens to magnify the solar image to fi ll most of your camera’s imaging sensor without too much loss of light.

In fact, ordinary digital cameras, whether compact or DSLR, are not ideal for H-alpha imaging. The reason is that in color cameras the pixels on the imaging sensors are covered by an array of tiny red, green and blue fi lters. This is fi ne in everyday photography and white-light solar photography, because we are shooting in all wavelengths of the visible spectrum. But because H-alpha light is a very pure red light, which includes absolutely no green or blue wavelengths, those pixels that happen to lie under the blue and green fi lters are unused. Therefore an H-alpha digital image is made up of just the red pixels, which form only a minority of all the pixels in the camera. Thus, for example, a 6-megapixel camera may effectively become only a 2-megapixel camera when used with an H-alpha fi lter. These 2 megapixels, furthermore, are spaced further apart than the original 6 megapixels. The result is that we get an image with, at most, one-third of the resolution and contrast that we would normally expect a 6-megapixel camera to deliver. The same is true for color webcams and dedicated astronomical CCD cameras of the “single shot” type – i.e., CCD cameras which download a color image after a single expo-sure, unlike monochrome cameras, which require you to take three separate images through different color fi lters and combine them in the computer later to create a color image. The very best H-alpha and Ca-K images that you see in astronomy magazines or on websites, showing intricate details and the fi laments and chromo-spheric granulation almost as dark as sunspots, were taken with high-end mono-chrome CCD and webcam-type cameras. The attractive colors of these images are, in fact, false colors introduced at the image processing stage.

Despite this limitation, however, you can take reasonable pictures with even an inexpensive compact camera and the afocal method (Fig. 7.15 ). One limitation with H-alpha photography is caused by the fact that prominences are generally much fainter than the solar disk and its features, so you can only capture prominences if you overexpose the disc. If you want a shot showing both the disc features and the surrounding prominences in the same image, you need to take two images, one exposed for the disc and the other for the prominences, and combine them in the computer afterwards. Making such composite images is very popular with amateur astronomers. However, if you aim a compact camera set to auto-exposure into a low-power eyepiece, showing the Sun surrounded by lots of black sky, you may well fi nd that your camera only slightly overexposes the disc, without completely washing out the fi laments and other disc details, while prominences are visible, at least faintly, on the same image.

The main trouble with H-alpha photography is focusing. If your camera is a compact, you will encounter the same problems with focusing with the LCD screen that are discussed earlier in the section on white-light imaging, and the problem is compounded by the fact that, unlike the white-light photosphere, the chromosphere has no hard edge on which to focus. One way round this is to “de-tune” your H-alpha fi lter so that its transmission is well off the H-alpha line and it shows just a red disc, with no fi laments or prominences. The somewhat fuzzy and ragged

170 7 Imaging the Sun with a Digital Camera

chromosphere at the limb of the Sun disappears and you can focus more easily on the sharp limb of the photosphere. Focusing the H-alpha image can be even more dif fi cult with a DSLR, because the H-alpha image shows up only faintly on the coarse-ground focusing screen. The screen can also cause the image to break up into strange patterns, making it hard to get a focus at all. You might fi nd that it is easier to get reasonable images with a cheap compact camera than it is with an expensive DSLR!

If you examine your digital H-alpha images carefully, you may fi nd that they are covered with a strange pattern of curved bands, which can in some cases detract from the quality of the image. This is known as “Newton’s rings” and is caused by the interference nature of the fi lter interacting with the pixelated imaging sensor in the camera. Not all combinations of camera and fi lter are affected by this. For example, in a Canon 300D, the effect shows up only very subtly on images taken at relatively high magni fi cation with the Coronado SolarMax 40 fi lter. If your images are affected, you could try experimenting with a different focal length; failing that, you may need to resort to a different camera.

Fig. 7.15 The Sun in H-alpha, taken by the author with a Canon PowerShot A640 hand-held to a 25 mm eyepiece on an 80 mm refractor equipped with a Coronado SolarMax 40 fi lter. This image demonstrates the dif fi culty in obtaining good H-alpha images with color cameras. Two fi laments, some plage detail and a sunspot are faintly visible

171Photographing the Chromosphere

If you have, or can gain access to, a coronagraph you can obtain spectacular images of prominences, because the Sun’s glare is then blocked by the occulting disc (Fig. 7.16 ). The Baader instrument, attached to an 80 mm refractor, gives an effective f/ratio of f/17 – quite fast compared to many H-alpha systems – allowing very short exposure times. Using 1/250 of a second at ISO 400 on a Canon 300D gives excellent results, although exposures sometimes need to be adjusted for especially bright or faint prominences. The effective focal length of the system is 1,365 mm, which gives a good disc size on the sensor and allows even large promi-nences to be accommodated on the frame. The coronagraph has a diaphragm for controlling the amount of scattered light in the system, and stopping this down about half-way gives the highest contrast.

The main problem with coronagraph photography is keeping the Sun hidden behind the disc. This demands accurate polar alignment, which is not easy in the daytime, when you cannot see Polaris or any other stars. Polar alignment is not a problem if your telescope is permanently set up, but if the instrument is portable you need to fi nd another way. One possibility is to align the mount in the normal way at night and make marks in the ground showing the correct positions for the tripod legs – assuming no one minds marks being made on the lawn! Use a trial and

Fig. 7.16 Prominences photographed through a coronagraph attached to an 80 mm refractor. A fi nely detailed prominence is visible to the upper right of this image originally shot on 35 mm fi lm

172 7 Imaging the Sun with a Digital Camera

error process, fi rst aligning the telescope roughly north–south using trees and other local landmarks and then adjusting the mount until the Sun no longer slips out from behind the disc.

If you have a sub-angstrom fi lter, you can, in fact, simulate the effect of a coronagraph by overexposing the image so that the prominences show up well and the disc is (probably) burned out. You can then black out the disc using an image processing program. This and other techniques are discussed in Chap. 8 .

173L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0_8, © Springer Science+Business Media New York 2012

Webcam Imaging and Image Processing

Chapter 8

If you are shooting the Sun at high magni fi cation, the sharpness of the image you get with a conventional digital camera (compact or DSLR) will be dependent on the steadiness of the seeing at the moment the shutter was released. Even in good see-ing, when the image is relatively steady visually, close-ups of sunspots and other solar features can be disappointing because they were taken when atmospheric turbulence affected the image. You can counter this problem, to some extent, by shooting a large number of images with your digital camera in the hope of getting at least one image taken in a brief instant of steady seeing. The large capacity of modern memory cards makes shooting large numbers of images eminently possi-ble, and indeed is an excellent way of getting good images at long focal lengths using eyepiece projection with a DSLR. But the most effective solution is to shoot a “movie,” perhaps a minute long and composed of hundreds or even thousands of digital images, and then select the sharpest images in the computer later, stacking as many of them as possible to create a super-sharp fi nal image. This can be achieved using a device that in recent years has caused a revolution in the quality of amateur images of the Sun, Moon and planets—the webcam.

Like conventional digital cameras, webcams were not originally designed for astronomy at all. Their purpose is to act as small television cameras that broadcast either live video footage or recorded movies over the Internet. For example, they allow people to see each other while talking to each other online, and for this pur-pose your laptop computer may already have a small built-in webcam. Astronomers use either a webcam proper, designed for everyday use and obtainable from a com-puter store, or a device built on the webcam principle and optimized for astronomical

174 8 Webcam Imaging and Image Processing

imaging. Here we can de fi ne the term “webcam” to mean any digital video camera that sends a digital movie to a computer. For astronomical applications, webcams come in three broad categories: everyday webcams; cameras made by telescope companies that are effectively ordinary webcams designed for use with a telescope; and specialized, monochrome (black-and-white) cameras used by more advanced amateurs.

This chapter discusses how to use a webcam to image the Sun, then how to pro-cess the resulting movies in your computer to obtain a super-sharp, high-resolution fi nal image. The fi nal section shows how you can enhance your solar images using image processing software, whether they were taken with a webcam or a conven-tional digital camera.

Webcams and Accessories

Do You Need a Webcam?

One disadvantage of webcams is that the image sensors in these devices are very small—typically only 6 mm (¼ in.) in their longest dimension, far smaller than the 23 mm of a typical DSLR. Therefore, in order to include the whole disc of the Sun, or even a substantial part of it, in the frame, you need to use a very short focal length, which results in an image with very low resolution in terms of arc seconds per pixel. If you want such whole-disc images, it is best to use a conventional digital camera. (One exception is if you want to do a live “webcast” of a special event, such as a solar eclipse, when you need the webcam’s movie capability and the reso-lution of the image is less critical.)

However, if we increase the focal length so that, say, a single sunspot group nearly fi lls the frame, the resolution of the image in arc seconds per pixel can approach, or even exceed, the theoretical resolution of the telescope. Atmospheric turbulence normally prevents us from achieving this, but in a webcam movie it is sometimes possible to capture a handful of frames taken in fl eeting moments of steady seeing that show detail down to the theoretical limit. Then, powerful image-processing software for use with webcams allows us to select the “good” frames in a movie with the highest resolution and then stack these frames on top of each other to assemble a fi nal image that has high quality as well as high resolution.

Advantages of Webcams

Webcams are generally very small and extremely lightweight, and so on even the smallest telescopes they cause no problems with balance—unlike some digital cameras. This lightness is especially advantageous if you want to do H-alpha

175Webcams and Accessories

(or Ca-K) imaging with a fi lter system that uses a star diagonal, since with such systems a conventional digital camera can be top heavy and swing around without warning, and could even fall out of the eyepiece receptacle.

Another advantage of webcams is that they do not require batteries or any power supply of their own, because power is provided through the USB (Universal Serial Bus) cable connecting the camera to the computer. The other side of the coin is that unlike digital cameras, webcams need a computer to function, so a webcam is less convenient if your equipment is portable and you need to set up a laptop as well as your telescope every time you want to image the Sun (Fig. 8.1 ).

You would probably get the most out of a webcam if your telescope is large by solar observing standards (say 100 mm or 4 in., or larger) and it is permanently set up in an observatory. Not only does this avoid the inconvenience of setting up every time, but it is with larger telescopes that webcams really come into their own in imaging the Sun, because with their ability to coax high-resolution images out of the unsteady daytime seeing, only webcams can exploit the full resolving power of a large telescope. Conventional “still” digital cameras only rarely achieve this. In any case, controlling the camera from a computer has its advantages. First, it is easier to focus the image—there are no problems with trying to focus using the fuzzy LCD screen of a compact camera or the coarse-ground focusing screen in a DSLR. Secondly, operating the camera by remote control means that there is no problem with camera shake caused by you touching the camera when you make exposures.

Fig. 8.1 The author’s 80 mm refractor with webcam connected to laptop computer

176 8 Webcam Imaging and Image Processing

Choosing a Webcam for Solar Imaging

If you already use a webcam to image the Moon and planets, you should be able to use it to some extent for solar imaging. Some webcams originally designed for every-day use have become very popular with astronomical imagers. One early example was the Philips TouCam Pro, later upgraded to the TouCam Pro II, which was used by many amateur astronomers worldwide to produce some remarkable lunar and planetary images. These models are sadly no longer in production, but you might pick one up on the second-hand market. It is worth shopping around online and investigat-ing astronomy Internet forums to look for currently popular models for astronomy. Note, though, that if you buy an everyday webcam you will need to buy or make a suitable adapter to mount it to a telescope, since webcams are not generally designed for astronomy. As with buying a compact digital camera, it is worth enquiring whether a suitable adapter is available for it before parting with your money.

If you are a serious solar observer and want to get on with taking good-quality images without worrying about how to attach the camera to the telescope, you are best off buying a webcam speci fi cally designed for astronomical imaging through a telescope. Meade and Celestron have each produced a webcam-type camera opti-mized for astronomical imaging. Note that neither camera comes with a lens, so they cannot be used for everyday purposes without modi fi cation. The Meade Lunar and Planetary Imager (LPI) is now no longer produced, but you may well fi nd it on the second-hand market. It cost around $125 when new. This camera has a built-in tube at the front made to fi t a standard 1¼-in. drawtube, enabling the camera to be mounted to the telescope as simply and easily as an ordinary eyepiece. Still cur-rently available is the Celestron equivalent, known as the NexImage (Fig. 8.2 ),

Fig. 8.2 Celestron NexImage astronomical webcam attached to the author’s 80 mm refractor

177Webcams and Accessories

price $230. Like the Meade LPI, this camera is very small and light and comes ready-made to mount onto a telescope with a 1¼-in. drawtube. As well as its own driver software, it also comes with the all-essential software required for selecting the sharp frames, and processing the image to produce the super-sharp fi nal picture—in this case a copy of RegiStax , the best-known webcam image processing package for amateur astronomers. (The NexImage comes with an older version of RegiStax —if you want the latest version, you will need to download it from the RegiStax website.) The image sensor has a resolution of 640 × 480 pixels—small by the standards of modern digital cameras, but exactly what is needed for imaging small objects at high magni fi cation. In fact, the sensor is the same as that used in the Phillips TouCam Pro II, so you have all the advantages of this highly prized astronomical imaging camera without the need to adapt it for mounting on the telescope or download the image processing software separately. I have owned a NexImage since 2007 and have been very pleased with its performance on the Sun and other Solar System objects.

The Celestron NexImage and similar color cameras are excellent beginners’ cameras and even for experienced observers may be all you need if you do pre-dominantly white-light imaging. However, as explained in Chap. 7 , color cameras are not ideal for imaging in H-alpha or Ca-K, because they use only one of the three sets of color pixels on the imaging sensor, effectively reducing the camera’s resolution to about one-third of its full resolution in megapixels. In recent years, this problem has been solved by the appearance of a number of monochrome CCD cameras which operate on the webcam principle—i.e., they produce live video output, from which sharp frames can be selected and stacked. These cameras are designed for scienti fi c imaging and some have been designed speci fi cally for astronomy. Although they are built to higher technical speci fi cations than begin-ners’ webcams, their prices are often surprisingly moderate, starting at around double the cost of the Celestron NexImage.

Some of the least expensive monochrome cameras are produced by the German fi rm The Imaging Source. For example, their DMK21AU04.AS monochrome CCD imager costs about $400. Its sensor is the same size as the NexImage—¼ in.—and it produces video output with the same resolution, 640 × 480 pixels. However, it can take up to 60 frames per second (fps), double the frame rate of the NexImage, so has the potential to capture more fl eeting instants of sharpness. The camera comes with its own image capture software. More “up-market” are the cameras made by the Canadian company Lumenera. For example, the monochrome version of their SKYnyx2-0 camera sells for $995. Its sensor is slightly larger in physical size than the Celestron and Imaging Source models, though it produces video output in the same format—640 × 480—and again shoots at up to 60 fps. Lumenera also produce higher-resolution cameras with larger imaging sensors, though at signi fi cantly higher cost.

More recently, some amateurs have begun using the tiny “Flea3” camera (Fig. 8.3 ) produced by another Canadian fi rm, Point Grey Research. This camera really is tiny, measuring barely 1¼ in. on a side, and is extremely lightweight. Yet the Flea3 can shoot at up to 76 fps and has a ½-in. sensor that gives 648 × 488-pixel video output—very similar to other cameras. This camera has larger pixels, which means, in theory,

178 8 Webcam Imaging and Image Processing

better-quality images, since smaller pixels produce more electronic noise. Note, though, that this model requires a FireWire computer connection—unlike the models by Celestron, The Imaging Source and Lumenera, which all use USB 2.0.

Computers and Accessories

Any computer that is not more than 4 or 5 years old and has a modern operating system (Windows XP or later if it is a PC) is capable of acquiring and processing webcam images. If your telescope is portable, you might want to consider using a computer of the “Netbook” variety—i.e., a very compact laptop without a CD/DVD drive. Netbooks are conveniently small and light and operate for a fair amount of time on battery power. Using a second computer like this also means that you don’t have to risk exposing your main laptop to the elements. Whatever computer you use, however, you will need plenty of disk space. The digital video fi les produced by webcams are huge by ordinary standards—typically hundreds of megabytes and potentially gigabytes in size. If you shoot lots of webcam videos you will quickly fi ll up your computer’s hard disk, which can slow down its performance. In the long run, consider using an external hard drive or USB memory sticks to archive your original webcam videos.

An important accessory for all webcam-type cameras is an infrared (IR) block-ing fi lter. The imaging sensors in these cameras have high sensitivity to infrared

Fig. 8.3 The tiny Flea3 webcam-type CCD camera on Dave Tyler’s 130 mm Astro Physics refractor mounted “piggyback” on his 14-in. Celestron SCT. The Flea3 is smaller than the Barlow lens it is attached to!

179Taking Webcam Images

radiation, and if your telescope is a refractor it will show up due to chromatic aberration. Even high-end apochromatic refractors are not corrected to eliminate infrared. Without an IR blocking fi lter, infrared shows up as a faint but distracting ring around the Sun’s limb and reduces the de fi nition of solar features. Infrared shows up most of all in classical achromatic refractors. Infrared should not, in theory, be a problem when imaging the Sun through solar fi lters, whether white-light or H-alpha, since they are designed to block infrared to enable safe visual use, but it is good practice to use a blocking fi lter, since it helps prevent dust from get-ting onto the image sensor. A blocking fi lter is, however, essential if you are using a Herschel wedge to image the Sun, as these devices transmit the full spectrum of the Sun’s light. IR blocking fi lters for webcams are inexpensive and screw into the front of a webcam like an eyepiece fi lter. Most webcams should have a nosepiece that allows the camera to be inserted into a 1¼-in. drawtube like an eyepiece, and this nosepiece should be threaded for standard 1¼-in. fi lters.

A motor-driven equatorial mount is not, in theory, essential for webcam imag-ing, as each frame of a webcam video is taken with a very short exposure, and the movement of the image due to Earth’s rotation can be corrected at the image pro-cessing stage. But a driven telescope is a requirement in practice, because the small image sensor of these cameras, combined with the large image scale at which we are shooting, means that the image moves out of view very quickly, and it is impossible to obtain a movie of decent length without some means of keeping the target in view. Also, electronic slow motions are pretty much essential for centering the target on the tiny sensor. However, the ability to correct for slight movement of the image during video capture means that precision polar alignment is not required for webcam imaging. Indeed, it is actually bene fi cial to have the image moving very slightly. Individual frames in a webcam movie have a rela-tively large amount of thermal noise, and so if every image of, say, a sunspot was aligned with a “hot” pixel, this pixel would show up as a distracting dot on the fi nal image. But if the image is moving relative to the sensor, any “hot” pixels will be in a different position on each frame and will be lost among the hundreds of frames being taken.

Taking Webcam Images

Good seeing is just as important for webcam imaging as it is with ordinary solar photography and visual observing. Although webcam imaging does much to coun-ter the effects of atmospheric turbulence, it is only in relatively good seeing that the brief moments of steadiness are frequent and long-lasting enough for a signi fi cant number of “good” frames to be built up to create a sharp fi nal image. In bad seeing there are few or no steady spells, and even with the best webcams and imaging techniques in the world the resulting images will be mediocre. You cannot process image data that is not there. The best webcam images are always taken under the best conditions. To get really sharp webcam images, showing details at or near the

180 8 Webcam Imaging and Image Processing

resolution limit of your telescope, you need to take the steps described in Chap. 3 to maximize the chances of good seeing.

High-resolution webcam imaging—indeed, any sort of imaging that pushes your telescope’s resolving power to its limit—demands more attention to detail than does imaging the whole solar disc or large parts of it. As well as the seeing, you need to consider the altitude of the Sun, especially if you are shooting in white light. As noted in Chap. 3 , if the Sun is very high in the sky, it may heat up the ground enough to cause a large amount of convection, which spoils the seeing. Conversely, if it is less than 30° high, atmospheric dispersion becomes signi fi cant, causing fringes of false color to appear in the image, and you may need to use one of the fi lters described in Chap. 3 to counter it. (Atmospheric dispersion is not a problem in H-alpha imaging, as here we are shooting at just one wavelength.) If your telescope is a re fl ector or catadioptric, it is very important to ensure that its optics are well-collimated. Imaging the Sun in white light with a refractor demands a judicious choice of secondary fi lters—that is, fi lters in addition to the aperture fi lter over the objective—to reduce chromatic aberration to a minimum. As well as the IR blocking fi lter, try using an interference fi lter such as the Baader Fringe Killer or Solar Continuum fi lter, described in Chap. 7 .

Because webcams usually have small image sensors and therefore small true fi elds of view, it may not be necessary to use a very long effective focal length to get a big enough image scale for your target sunspot or prominence to have an appreciable size on the image and to get the most out of your telescope’s resolution. As noted above, although using too short a focal length will not exploit the tele-scope’s full resolution, using too large an image scale will reveal no additional detail and will only make the image fainter and increase the effects of atmospheric turbulence. Therefore, you may fi nd that it is unnecessary to use eyepiece projec-tion to achieve the extreme focal ratios that you would when using a DSLR. A Barlow lens may be all you need, even with a small telescope. Make sure the Barlow you use is a good-quality one, as some cheap Barlows cause chromatic aberration and internal re fl ections. You can use a Tele Vue 2.5× Powermate for all your webcam imaging of the Sun through a 60 mm Takahashi refractor, which has a focal length of only 355 mm when used at prime focus.

The procedure described here for taking webcam movies is based on the author’s own experience with a Celestron NexImage camera and its image acquisition soft-ware, AMCap. But as webcams all work on the same principle, it should not be dif fi cult to adapt the procedure to suit your camera and software.

When you have pointed your suitably fi ltered telescope at the Sun, and installed an appropriate Barlow lens or other arrangement to magnify the prime focus image, begin by plugging the camera’s cable into your computer’s USB or FireWire port, and then start up the camera’s image capture software (which should be installed on your computer’s hard drive before you begin imaging). Under the “Options” menu, select “Preview” if this option is not already ticked. Now remove the cap from the camera’s nosepiece, taking care not to point the camera at the Sun or any other bright light source when it is uncovered. You should see the image capture program window light up as the image from the camera is displayed.

181Taking Webcam Images

Carefully insert the camera into your telescope’s drawtube, ensure that it is fi rmly clamped and then focus the image. You should now see a live TV picture of the Sun on your monitor, though how much detail you see depends on the camera’s settings. Go back to “Options” and select “Video Capture Filter.” This will bring up a dialogue box in which you can adjust several settings. Adjust the “Brightness” and “Gain” until a recognizable image of the Sun is displayed. Brightness adjusts the brightness of the image as seen on your monitor, while gain varies the sensitiv-ity of the camera—i.e., it effectively sets the ISO rating of the camera. As with a high ISO setting on a still digital camera, setting the gain to maximum will give remarkable sensitivity but a heavily speckled, “noisy” image. Turning the gain down to minimum gives a fainter but much smoother image. As the Sun is a bright object, even through a solar fi lter, only a low to medium gain should be necessary for recording its image on a webcam. “Video Capture Filter” also allows you to set the shutter speed. As you become more experienced with webcam imaging, you will want to try different exposure times, but as a fi rst experiment try setting the exposure time to “Auto.” If your target is a sunspot or group of sunspots near the limb, move the telescope as close to the center of the Sun’s disc as possible while still keeping the target sunspot in view, so as to minimize the amount of black sky in the image, which can cause the camera’s auto-exposure to overexpose the image. Click “Apply,” then “OK” to exit “Video Capture Filter.”

With the correct brightness and gain settings selected, the image of the Sun should now be realistically displayed on your monitor, much as you would see it in the telescope eyepiece, complete with the fl ickering and “boiling” of the image due to atmospheric turbulence. Now is the time to focus the image more precisely. Good focusing is at least as important with a webcam as it is with a conventional camera, as you will not be able to exploit the spectacular resolution of webcam images if they are even slightly out of focus. Night sky imagers use some elaborate tech-niques for precision focusing, such as knife-edge focusing or Hartmann masking, but these work using star images, which are impossible during daytime. Fortunately, if you are shooting in white light, the Sun’s limb provides a hard edge on which to focus. Focusing is slightly more complicated in H-alpha, since the limb is sur-rounded by the chromosphere, whose outer edge is ill-de fi ned due to the spicules (see Chap. 6 ). The best way around this is to “de-tune” your H-alpha fi lter so that its peak transmission is well off the H-alpha line. This causes the chromosphere, as well as any other H-alpha features such as prominences, to disappear, and the Sun simply appears as a red disc with a well-de fi ned limb, which is easier to focus on. When the “de-tuned” solar image is in focus, you can then re-center your fi lter on the H-alpha line.

You then need to select the resolution of the video image, the compression and the frame rate. Resolution and compression should not be changed after they are set up for the fi rst time. Resolution should always be set to 640 × 480 (VGA). If your camera offers a higher resolution, it will likely achieve this using interpolation, a clever technique that simulates higher resolution but does not truly increase the resolution, and should not be selected. Always set compression to “None” if your camera offers this; otherwise, select “I420.” As in a conventional camera, the

182 8 Webcam Imaging and Image Processing

greater the image compression, the more data is lost and so the less is the image depth. An advantage of the more expensive cameras over the cheaper ones is their ability to shoot uncompressed video. You should normally set the frame rate to the maximum offered by your camera, as this gives you the greatest chance of captur-ing sharp images.

Before you begin imaging, you need to set the “time limit” for the video you are about to shoot—i.e., the total exposure time of all the frames. This is found in the “Capture” menu in the NexImage software. The number of frames in a webcam video is equal to the time limit multiplied by the frame rate, so if we shoot a video lasting 30 s (i.e., with a 30-s time limit) at 30 fps, our resulting movie will be com-posed of 900 frames—enough to occupy nearly half a gigabyte on your computer’s hard drive. Longer movies result in even larger fi les, as do movies of the same length shot at a higher frame rate. Finally, you need to set a fi le name to save the movie to (with the NexImage, choose “File,” then “Set Capture File…”). With NexImage it is important to always remember to do this before you start shooting, as otherwise the system automatically saves the latest movie to the fi le name of the last one it shot, and so overwrites your previous effort. Save each video with a recognizable fi lename that includes the date, time and perhaps also the NOAA number of the active region you are shooting, if you know it in advance.

You are now ready to begin shooting. With the NexImage, select “Capture,” then “Start Capture.” Shooting will stop automatically at the end of your chosen time limit. Start by taking a number of short movies, perhaps of no more than 100 frames each. This will give you fi les of manageable size and will allow you to experiment with the focus and camera settings without taking up too much disk space. When you play back a webcam movie, you will see the image moving about slightly and rippling in and out of focus with the seeing, like the view through an eyepiece.

It is now time to bring the sharp frames out of the movie and use them to the best advantage to assemble the fi nal image.

Processing Webcam Images

If you play back a movie of the Sun shot through a webcam, each individual frame will not look particularly exciting, even sharp frames taken in moments of good see-ing. The image will appear dull and somewhat “grainy,” with rather low resolution—lower than you would expect in an image taken with a conventional digital camera (Fig. 8.4 ). This is partly because the image sensor in a webcam has much higher electronic noise than that in the best digital cameras, and also because the resolution of a webcam, in terms of pixels, really is low—little more than 0.3 megapixels. But if we stack a large number of webcam images on top of each other and merge them to create a single fi nal image, the noise becomes insigni fi cant and the image takes on a smooth appearance like an ordinary digital camera image. And because we are shooting at a very large image scale, with our target solar feature fi lling a good proportion of the frame, the puny 0.3-megapixel resolution does not matter.

183Processing Webcam Images

Webcam image processing works in three stages. First, the images are aligned around one or more reference points in the image, such as a small sunspot. Then the images are checked for quality and the poorest (i.e., the most blurred) ones discarded. Next the aligned images are stacked on top of each other and merged into a single image fi le. This greatly increases the contrast of the combined image. Finally, the image is sharpened using a feature called “wavelets.” It is this fi nal stage that gives webcam images their incredible sharpness and detail.

Remarkably, all this can be achieved using software packages that can be down-loaded from the Internet free of charge. The best-known webcam image processing program, and the best-loved among amateur astronomers, is RegiStax ( www.astronomie.be/registax ). This comes free with the Celestron NexImage camera, but if you use other cameras, or you use the NexImage and want the latest version of RegiStax , you need to download it from the website. Other webcam image process-ing programs include Astrostack ( www.astrostack.com ) and AviStack ( www.avistack.de ).

It is essential to remember two important points about processing webcam images. First, although image processing programs can work wonders with what appears to be a mediocre set of images, you need to be very wary of over-processing your images. Solar features should look natural in the fi nal image, as you would see them visually at high magni fi cation in good seeing. If solar features have very hard edges, or if sunspots have white rings around them, the image is over-processed.

A common feature of over-processed images is the presence of masses of what looks like extremely fi ne detail, such as tiny pores or details within the granulation, but which is actually image processing artifacts. Related to this is the fact that, as

Fig. 8.4 A single frame from a webcam movie has a grainy appearance and low resolution. This image of a group of prominences was taken by the author on February 11, 2008. It is one frame from a 900-frame movie taken with a Celestron NexImage webcam attached to a 60 mm refractor with a Coronado SolarMax 40 H-alpha fi lter

184 8 Webcam Imaging and Image Processing

mentioned above under “Taking Webcam Images,” to capture the fi nest detail with a webcam you still need good seeing. No amount of processing will turn a webcam movie shot in poor seeing into a perfect fi nal image. Although it is always worth trying to see what detail you can extract from a movie shot in such conditions, there will be a limit to the amount of detail you can resolve and all you will get from excessive processing is image artifacts. Remember: you cannot enhance detail that is not there.

Here we shall brie fl y discuss how to process a solar webcam image using RegiStax . When you fi rst start up RegiStax , it may look complicated and intimidat-ing, but the basic principles of its use are actually quite simple and you can gradually learn the more advanced features as you become more experienced. The program works through various screens, marked by a series of tabs at the top of the applica-tion window. When you open RegiStax it should display the fi rst of these, marked “Input.” Begin by clicking “Select Input,” which brings up a standard dialogue box enabling you to select the webcam movie you have taken from the directory you have saved it to. Webcam movies will have the fi lename extension “.avi.” Opening the movie will bring up a still image, the fi rst of however many frames you shot. Tick the box marked “Show frame list,” and this will bring up a list of all the frames. Clicking on each one on the list will change the preview window to the frame selected, e.g., clicking on “Frame 10” will show the tenth frame in your movie.

To align the images, you need to select a sharp frame from this list to act as a reference frame, with which all the other selected frames will be aligned. Go through the images one by one and select which one you think is the sharpest. This is another reason for shooting a short movie to begin with, as it will take a long time to go through 900 frames if you run the camera for 30 s at 30 frames per second! Alternatively, click “Set None” at the top of the frame list. This will de-select all the frames, and you can then tick the box at the side of each frame on the list to select, say, the fi rst 150 frames. Next, choose an object in your reference frame that you can use as your reference point for alignment—a small, compact sunspot is ideal, as is a small prominence or a bright spot within a prominence. When you move the mouse pointer over the preview window, it should change into a small box shape. You need to adjust the size of this box so that it is big enough to neatly sur-round your chosen feature—select from various sizes in the list marked “Alignment box,” then click the mouse when the box is over the object. (The latest version of RegiStax allows you to use more than one alignment reference point.)

RegiStax should now advance automatically to the “Alignment” tab. Without going into the technical details here, prominent in this new screen are two windows: on the left is a multi-colored map of the image showing how accurately aligned the various frames are on top of each other, while to its right is a graph showing various settings for the quality of the image. The alignment map will have a red area at the center of the image; when this is as small as possible, the alignment of the frames is good. Clicking “Recalc” to the upper right of the screen improves the alignment and reduces the size of the red area. Leave the image quality settings to their default values until you gain more experience, as these are quite complex, and even the default values will achieve remarkable things with an average-quality image.

185Processing Webcam Images

Finally, click “Align & Stack” at the bottom right of the screen. This command both throws out the most blurred images and then aligns and stacks the remaining ones. This process may take a few minutes, depending on the speed of your com-puter and the size of your original video fi le. If you have chosen a good reference frame, an appropriate reference point on the image and the correct alignment box, you may be pleasantly surprised at the result: a bright, high-contrast and, above all, smooth combined image, a huge difference from the dull, grainy appearance of the individual video frames (Fig. 8.5 ).

The image is then sharpened using the tools provided under the fourth tab, called “Wavelet Processing” (Fig. 8.6 ). Wavelet processing splits the image into six lay-ers, each of which can be adjusted separately to bring out details in the image. Each layer is adjusted using sliders on the right-hand side of the screen, and you see the combined result of adjusting the sliders on the image window to the left. Layer 1 adjusts the coarsest details in the image, layer 2 looks after somewhat fi ner detail, and so on down to layer 6, which handles the smallest details of all.

Generally, with an average image only layers 1, 2 and 3 make much difference to the image, with layer 1 giving the most dramatic changes. Be very careful when using wavelet processing: while it can dramatically improve an image, it is also at this stage that it is easiest to let your enthusiasm for image processing run away with you and “overcook” the image. Only move the sliders far enough to improve the image while keeping its appearance natural. When the solar features start to have unnaturally hard edges and exaggerated contrast, it is time to stop. Moving sliders 1 and 2 to around 20 or 30 points out of 100 will generally enhance the image without exaggerating it too much. It is a good idea to experiment with different wavelet settings and save the image each time, making notes on the

Fig. 8.5 The same prominences as in Fig. 8.4 , after the webcam movie has been through the stacking and aligning process

186 8 Webcam Imaging and Image Processing

settings you use. A quick way of noting the settings is to save the images with fi lenames that include what has been done to the image—e.g., “Sun 20080211 21,29,11.”

Enhancing Digital Images

A solar image taken with a webcam and processed using RegiStax or a similar program may be ready for viewing and may not need further enhancement, since it has already been subjected to considerable processing to enhance the contrast and bring out detail. But when you fi rst see images from a conventional compact digital camera or DSLR on your computer screen, they may appear disappointing, with sunspots looking washed out and the whole picture somewhat fuzzy. But don’t be discouraged, as such images contain a lot of data that you can extract using the computer. This is where image enhancement using image processing software comes in. Some image enhancement techniques, such as adjusting color, are also useful for webcam images.

Digital techniques can also be useful if you took many solar images in the age of fi lm. Digital solar imaging only completely took over from fi lm photography in the 2003–2005 period, so your images of the great sunspot groups and prominences from the 2000 solar maximum might well be on fi lm. You can scan conventional slides, negatives or prints into computer-readable format and then enhance them with image processing software in the same way as images taken with a digital sys-tem. Additionally, a scanned digital image may also be the only way a laboratory can make a good print from a slide or digital image nowadays, in an age when chemical photographic laboratories are fast disappearing.

Fig. 8.6 The fi nal image of the prominences, after wavelet processing

187Enhancing Digital Images

For scanning your old photographic prints an ordinary A4 fl atbed scanner, avail-able very cheaply from any computer store, is fi ne. You might well have a suitable scanner as part of a 3-in-1 printer-scanner copier. To scan slides or negatives you need a fi lm scanner. These are more expensive, but much cheaper than they used to be. An important point to bear in mind when scanning images is always to scan the image in TIFF (.tif) or bitmap (.bmp) format. If at all possible, avoid using JPEG (.jpg) format; as explained in Chap. 7 , JPEG compresses the image data—that is, it leaves information out of the image fi le to keep the fi le size small—at the expense of image depth, causing what should be a smooth gradation of tone to show up as an unsightly array of layers. Scanning in uncompressed format gives the image a large fi le size on your computer: a bitmap image with a resolution of 1024 × 768 pixels (ideal for displaying on a computer screen or with a digital projector, e.g., for showing solar images at astronomy meetings) has a fi le size of about 2.5 MB, while a full TIFF scan can take over 20 MB of disk space. But at least you can be sure that you have transferred all the data from the original photographic image. You can always convert your image to JPEG later if you wish to send it by e-mail or post it on a website; how to do this while avoiding too much compression is explained later in this section.

Software for processing images varies greatly, both in price and the number of features it contains. The brand-name image processing packages used by profes-sional photographers and publishers are very expensive. For example, the current version of Adobe Photoshop (CS5) retails for £650 (though it is now possible to spread the cost by taking out a monthly subscription to this software). But many programs containing all the features you need for processing solar images are avail-able for much less than this. A really good image processing package is a cheaper version of Photoshop called Photoshop Elements , which can be bought from cam-era stores or online, but often comes “bundled” with a new digital camera, printer or scanner. Indeed, if you are choosing a new digital camera, it might be worth investigating if any image processing software comes free with it. Elements has many of the features of the full Photoshop and more than you will ever need for solar photography.

An alternative to Elements is Corel’s Paint Shop Pro , which sells for a similar price. You might also wish to try a program speci fi cally designed for astronomical imaging, such as AIP4Win , which comes with Berry’s and Burnell’s Handbook of Astronomical Image Processing , or MaxIm DSLR, which offers remote control of many DSLR models as well as image processing. RegiStax offers some basic image manipulation functions, such as cropping and color adjustment, once the image has been through wavelet processing, but for more advanced enhancement you will still need an image processing package.

Before beginning to process images, it is essential to remember two golden rules. First, when you have chosen an image you want to enhance, always work with a copy of the fi le, never the original image. This precaution ensures that you can always return to the original image if things go wrong during processing. It is not safe to rely on the “Save Changes?” dialogue box when you close an image, as it is only too easy to click on “Yes” by mistake and so lose your original image forever. When you download the original images from the camera, save them with

188 8 Webcam Imaging and Image Processing

the date in the fi lename—e.g. “201103080004” would be the fourth image you took on March 8, 2011. Then, use “Save As” to save the copy with an extension indicat-ing what intend to do with the image: for example, “201103080004en” would mean an enhanced version of 201103080004.

The second golden rule is never to “overcook” your images through too much processing. This applies to images from digital cameras or fi lm as much as to web-cam images. In particular, enhancing brightness and contrast too much will produce artifacts such as bright rings around sunspots. If working with images originally shot on fi lm, bear in mind that fi lm has a limited dynamic range and so there will be a limit to how much you can enhance your images. Scanning slides is a good way of viewing the slides easily on the computer screen, and by scanning either slides or negatives you can be surprised how good the original image was, espe-cially if you are familiar only with the photographic print made from a slide or negative. Photographic laboratories were never geared towards astronomical imag-ing, and so often produced disappointing prints of solar photographs. When you see a scanned image from the original negative, the improvement can be dramatic.

To demonstrate how the various features of an image processing program can improve a solar image, let us use Photoshop Elements to enhance one of the author’s solar images, taken on March 8, 2011, with a Canon 300D on an 80 mm refractor with a 1.4× teleconverter and a Baader AstroSolar Photo Film fi lter (Fig. 8.7a ). Other image processing programs perform the basic operations described here, so you should be able to follow the procedure with your own software.

File Formats

To shoot an image with no data compression, you need to shoot in the camera manu-facturer’s proprietary “RAW” format. Unfortunately, most common image process-ing programs cannot read this format directly, and most DSLRs do not give you the option of shooting directly in TIFF. Before working with a RAW image, you need to convert it to TIFF using the software provided with the camera for downloading images. This can make the fi le very large on your computer; for example, a RAW image from a Canon 300D takes up about 6 MB, but converting it to TIFF increases it to 36 MB! Fortunately, modern computers can easily handle fi les as large as this, but hard disk space will be an issue if you take lots of images. In any case, it is good practice to back up images regularly onto some removable medium, such as a USB memory stick or CD-ROM, in case your computer’s hard disk stops working or the computer itself is lost, damaged or stolen, as can happen if it is a laptop.

If the original image is a JPEG (as it probably will be if it is from a compact digital camera), take care to save it so that there is no further compression and hence loss of image data. Some programs default to a medium quality setting when you save an image, which can compress the image by a further 50% or more, and so when you save a copy of your image you can lose image quality without know-ing it before you even start working on it. In Elements , under “JPEG Options,”

189Enhancing Digital Images

choose a quality setting of 12 (maximum). This will give the largest possible fi le size and will minimize any further compression.

When you have fi nished working with an image and want a version with a smaller size—for example, for sending by e-mail or posting on a web page—you

Fig. 8.7 Steps in enhancing a digital image of the Sun using Adobe Photoshop Elements . ( a ) Original image, taken by the author on March 8, 2011, with a Canon 300D on an 80 mm refractor with a 1.4× teleconverter and a Baader AstroSolar Photo Film fi lter, ( b ) The image cropped to center the Sun’s image and exclude some of the black background, ( c ) The image fl ipped to show east to the right, ( d ) The brightness decreased by 20 points and contrast increased by 40 points, ( e ) Image after applying “Sharpen” twice and then with contrast and sharpness enhanced further using “Unsharp Mask,” ( f ) Image with red and green channels adjusted to give a distinct yellow color

190 8 Webcam Imaging and Image Processing

should choose a moderately high quality setting, such as 9 or 10 out of 12. This gives compression that is usually too moderate to be noticeable and it reduces the size of the fi le by a surprising amount. Elements will display the fi le size as you move the quality slider bar between “small fi le” and “large fi le.”

Cropping and Trimming

In Fig. 8.7a the original image looks washed out and with the sunspots lacking contrast. The fi rst thing we can do to make the image more aesthetically pleasing is to trim the edges of the photograph to show the solar disc in the center of the picture and fi lling most of the frame. Selecting the “crop” tool in Elements causes the mouse pointer to turn into this program’s crop symbol. Position the pointer towards the top left corner of the picture, then hold down the left mouse button and “drag” the mouse down and to the right. The selected area of the image appears as a rectangle outlined by dotted lines, whose shape varies as you move the pointer; dragging it to the right gives a horizontal rectangle, while dragging it down increases the vertical height of the image.

When you have selected the area you want to include in your fi nal image, release the mouse button. Now click the right-hand mouse button and select “Crop.” The outlying parts of the image then disappear, leaving just the selected area. Figure 8.7b shows the result of cropping our example image. Of course, cropping an image enlarges it slightly, increasing the size of the pixels relative to the whole picture, but provided you do not crop the image too heavily—for example, by excluding everything except one small section of the picture—the pixels should not be notice-able. How heavily you can crop an image and still get away with it depends on the size of your camera’s image sensor in megapixels, and this is one advantage of a camera with a high megapixel count. Cropping also allows you to exclude parts of an image that are blurred or otherwise sub-standard and detract from the quality of the picture. This is especially useful if your camera is mounted separately from the telescope and part of the Sun’s image is missing.

Image Orientation

An image taken with a digital camera using the afocal method has the same orienta-tion as the visual view through the eyepiece. However, this orientation may be dif-ferent from what you are used to seeing, especially if you normally use the projection method, which reverses the east-west orientation. If you shoot through a Newtonian or a refractor without a star diagonal, north and south will be reversed as well. Turning the eyepiece around to a horizontal position in a refractor or catadioptric (see Chaps. 3 and 7 ) also changes the orientation. Image processing allows you to quickly and easily turn the image around so that it matches drawings and other photographs. Amateur drawings use a standard orientation showing north at the top

191Enhancing Digital Images

and east to the right. If you send images to a solar observing organization they should have this orientation.

To change the north-south and/or east-west orientation of your image, open the “Image” menu, select “Rotate,” then click “Flip Vertical” to fl ip north-south or “Flip Horizontal” to fl ip east-west. Figure 8.7c shows the example image fl ipped horizontally to show east to the right.

Changing the Brightness and Contrast

How well the sunspots and other features show up on any image depends on the contrast. Many image processing packages have a feature called “Quick Fix” or “Instant Fix,” which adjusts the brightness and contrast of an image automatically. In Elements , select the “Enhance” menu then click “Quick Fix…”. Accept the default settings of “Brightness” and “Auto Contrast.” Although this feature cer-tainly enhances the visual impact of an ordinary photograph, it is often not adequate for solar images, which need more adjustment. It may be suf fi cient if your image already shows solar details with good contrast but is a little too dark.

Most solar images require use of the speci fi c “brightness and contrast” feature. In most image processing packages this takes the form of a box with individual scales for brightness and contrast. In Elements , select “Enhance,” then “Adjust Brightness/Contrast,” then “Brightness/Contrast.” This brings up a box with individual scales for brightness and contrast. You can set the brightness and contrast either by entering numbers between these values in the boxes provided or by sliding a pointer along each scale using the mouse. Make sure that the “Preview” option is ticked, so you can see the effect of any changes “live” without having to close down the menu item.

It is interesting to play around with the brightness and contrast scales to see the effect of changing the values. Note that increasing the contrast value alone does not improve the contrast of the sunspots; in fact, it makes the image look brighter and reduces the contrast of the spots. To increase the contrast, fi rst decrease the bright-ness value, then increase the contrast by somewhat more. The spots, limb darkening and any faculae near the limbs will now show up much more prominently, while the photosphere will still be surprisingly bright.

If your original image is overexposed, you must decrease the brightness value substantially to make the spots show up better. If your camera has full manual exposure, your image will likely have approximately the correct exposure, as you will probably have experimented with a number of test exposures and previewed them on the camera’s LCD screen. Such an image will require only moderate brightness and contrast adjustment. But it is surprising how much even a correctly exposed picture can be “livened up” by adjusting the brightness and contrast. In Fig. 8.7d the brightness value has been decreased by 20 points out of a possible 100 and the contrast increased by 40 points. The sunspots are now much darker, and the image now shows several smaller spots that were not visible on earlier versions.

Also, bear in mind that an image can look misleadingly bright when you view it on the camera screen, especially if you have turned up the brightness to improve

192 8 Webcam Imaging and Image Processing

the visibility of the screen in daylight. When you download the image to your com-puter, you may fi nd that what you thought was a bright enough image could have done with a slightly longer exposure.

Removing Dust and Scratches

As noted in Chap. 7 , dust can be a major problem in solar photography. If you use a compact digital camera to photograph the Sun, dust particles on the telescope eyepiece can show up as unsightly gray spots on the image, while if you shoot with a DSLR, dust on the imaging sensor appears as black specks that could be mistaken for sunspots. Dust on 35 mm slides will also show up when you scan them. In theory, digital processing allows us to remove these. Some image processing pro-grams have a feature whereby you can erase dust spots above a speci fi ed size and intensity. But using this can cause particular problems for solar imagers. It removes the dust particles all right, but it can also remove all the sunspots!

Don’t make the mistake of thinking that a software feature called “Dust Brush” literally means a brush inside the scanner that cleans the slides. On using it with a slide containing a solar image, you might be horri fi ed to fi nd that the software removes not only the dust but also most of the sunspots as well! Dust removal com-mands also have the effect of blurring the image. With care this feature can give some improvement, but on the whole it is not recommended that you use it.

The best way of dealing with dust in solar images is to use clean optics, imaging sensors and fi lm originals to begin with. But a small number of dust particles can still get into an image and, if they are not tolerable, you can “paint” them out using the “Clone Stamp” tool in Elements . Select an area of the photosphere that is of the same brightness as that affected by dust, using the tool bar to vary the size of the “stamp” so that it is slightly larger than your target dust speck. Use the stamp to copy this “clean” area of photosphere onto the dust spot, which should then disap-pear, or fade to near-invisibility. You must get the brightness and color exactly right, as even a slight error can give a ghastly blotch on your image, must worse than the original dust spot. (This demonstrates the importance of only working with a copy of the original image). In any case, always note carefully where you have altered parts of the image and with what technique, as an image that has been altered in parts is no longer scienti fi cally accurate.

Sharpening the Image

Our image (Fig. 8.7d ) now looks far better than the original, but it is still not quite as sharp as it could be. You can correct minor blurring such as this at the click of a mouse by using the “Sharpen” feature (under “Filter” in Elements ). Using “Sharpen” once or twice can make details in the picture look crisper. Sharpening

193Enhancing Digital Images

does not increase the resolution of the image, it just makes the detail stand out better—you can’t resolve detail that is not already there! Don’t apply “sharpen” too many times to one picture, as this makes the image more “noisy,” detracting from the overall quality of the picture.

It is possible, however, to extract even more information from a digital image using an image processing feature known as “unsharp mask.” This allows you to enhance the contrast and sharpness of the image further while controlling the amount of noise in the picture. Selecting “Unsharp Mask” in your image processing package (“Filter,” “Sharpen,” “Unsharp Mask…” in Elements ) brings up a box with three sliding scales named “Amount,” “Radius” and “Threshold.” “Amount” runs from 1 to several hundred percent, “Radius” from 0.1 to 250 pixels and “Threshold” from 0 to 255 levels. Setting the amount to a high level makes the picture look sharper, but at the expense of excessive image noise. Radius controls the contrast: selecting a high radius makes the image look bright and stark, especially towards the limb, while reducing it lowers the contrast. Threshold makes the image smoother (i.e., reduces the amount of noise) when set to a high level but lowers the contrast, while a low threshold gives a noisy picture and makes the contrast too high, enhancing every dust particle and showing the sunspots surrounded by bright rings, which could be mistaken for faculae!

Clearly, we need to fi nd a compromise between contrast, sharpness and noise using the three controls. You might get your best results by setting the amount to a medium value (around 200%), the radius to between 3 and 20 pixels and the thresh-old to a fairly low value, around 25 levels. Figure 8.7e shows our example image, which has already been through “Sharpen” twice and has had its brightness and contrast adjusted as described above, treated with the unsharp mask feature with the amount set to 200%, the radius at 5.6 pixels and using a threshold of 25 levels. The result has even higher contrast, and more intricate sunspot detail has come “out of the woodwork” but without increasing the noise too much or causing facula-like rings around the spots.

Changing the Color

As noted in previous chapters, the color of a solar image is not important, as the Sun is really a monochrome object. However, it is sometimes useful to be able to change the color of an image to make it more aesthetically pleasing, as some fi lters give the Sun a lurid color. Some Mylar fi lters, for example, give the Sun a strong blue tint, while many glass fi lters make it look unnaturally orange.

With an image processing package it is easy to change the color of an image by increasing and decreasing the amounts of various colors present in the image. In Photoshop or Photoshop Elements you can do this professionally by splitting the image into three color “channels” representing the primary colors: red, green and blue. In Elements , select “Enhance,” “Adjust Brightness/Contrast,” then “Levels…”. From the drop-down list labeled “Channels,” select one of the primary colors, then

194 8 Webcam Imaging and Image Processing

move the sliders under the “Input Levels” histogram until you achieve the desired color, leaving the “Output Levels” sliders unchanged. You may have to adjust more than one color channel to give the image the correct tint; for example, to enhance the yellow color in Fig. 8.7f you can adjust both red and green. Some amateur astronomers have used “Levels” to enhance their H-alpha images so that a single exposure taken with a digital camera shows prominences and disc detail, because while H-alpha light, by de fi nition, should only go into the red channel, some cam-eras suffer “leakage” of red light into green and blue channels.

A quicker, though less precise way of changing the color of an image in Elements is known as “Hue/Saturation…,” found under “Enhance,” then “Adjust Color.” This allows you to vary the yellows, cyans and magentas as well as the reds, greens and blues. Select the color you want to change, then slide the “Saturation” adjustment along the scale between −100 and +100, as you would for brightness or contrast. This allows you to vary the yellows, cyans and magentas as well as the reds, greens and blues. Decreasing the blue, for example, is useful for reducing or eliminating the blue tint of a Mylar image, while increasing the blue gives a more pleasing color if your fi lter gives a strongly orange image. Decreasing the red is useful for correct-ing any reddening of the image caused by the Sun’s low altitude. Reducing the color saturation by a smaller amount is a good idea if you are happy with the actual color of your original image but would like it to be less intense.

Note that changing the color from, say, blue to yellow is not a substitute for using a secondary fi lter such as a light yellow or Fringe Killer fi lter. Secondary fi lters work by changing the wavelengths of light transmitted to the camera and so reduce chromatic aberration in refractors. Changing the color of the image on the computer, however, will make no difference to the effect of chromatic aberration.

Reducing the saturation by the maximum amount effectively turns your image into a black and white picture. Some programs also have a “remove colour” feature that allows you to turn the image to monochrome at the click of a mouse (in Elements , choose “Enhance”, “Adjust Color”, then “Remove Color…”)—though note that this is not the same as a true black and white image, which can be achieved by turning the image into a “grayscale” image (“Image,” “Mode,” “Grayscale”). 1

Conversely, if your original image is in monochrome—such as a solar image from years ago that you shot on black and white fi lm, or an H-alpha image taken on a monochrome webcam—it is easy to add color to give it an appropriate color, such as yellow for a white light image or red for an H-alpha one. Again, with Photoshop or Photoshop Elements you can do this using “Levels” and adjusting the red, green and blue channels. Alternatively, in Elements you can add color in a more roundabout way by selecting “Enhance,” “Adjust Color,” then “Color Variations…”. This brings up a series of “buttons” showing little thumbnails of

1 Some digital cameras also offer a black and white mode, though again it is not true black and white, because the image is still shot through the red, green and blue fi lter system on the imaging sensor. The only true monochrome cameras are dedicated astronomical cameras, such as the CCD cameras for night sky imaging or the webcam-type video cameras discussed in this chapter.

195Enhancing Digital Images

your image with a variety of tints. Click on these, or a combination of these, to achieve the color you want. Figure 8.8 shows a whole-disc solar image, originally shot in 2001 on Kodak Technical Pan 2415 black-and-white fi lm (now no longer made), and more recently scanned and tinted yellow using this technique.

Such “false color” images of the Sun can make your pictures more eye-catching and so are useful if you want to try your hand at having your images published in magazines or want to post them on websites. But the colors have no scienti fi c basis, and you should always state when false color has been applied. In recent years it has become fashionable to tint H-alpha images with a rich yellow, which gives a misleading impression as to the true color of an H-alpha image through the telescope.

Making Composite Images

A common problem encountered when imaging in H-alpha is that the prominences are much fainter than the disc, so that an image correctly exposed for the disc will show the prominences only faintly, if at all. Conversely, you may have an image on which the disc is overexposed but the prominences show up well. Many amateur astronomers get round this problem by creating a composite image composed of two photographs stacked on top of each other—the short exposure for the disc and

Fig. 8.8 Whole-disc solar image, originally taken by the author on June 19, 2001, on Kodak Technical Pan 2415 black-and-white fi lm, then more recently scanned and tinted yellow using Adobe Photoshop Elements

196 8 Webcam Imaging and Image Processing

the longer exposure for the prominences. The result is a rather beautiful image showing the disc surrounded by bright prominences, with the latter being even more easy to see than they are visually through the telescope.

Begin by opening both the image that shows the disc features at the correct exposure, and a copy of the overexposed image with the prominences. In Photoshop or Elements , select the “Magic Wand” tool, then click on the black background of the disc features image. This selects the background for cutting, copying or pasting, but by clicking “Inverse” under the “Select” menu, this will change the selection to the solar disc, which you can then copy to the clipboard as you would copy a selec-tion of text in a word processing package. Then switch to the overexposed image, and use “Edit” and “Paste” to paste your correctly exposed solar disc onto this image. When you have done this, the two discs will probably not be aligned, but selecting the “Move Tool” on the palette allows you to move it to the correct posi-tion by pushing on the mouse. Pasting one image onto another like this creates a second layer to the image, so you fi nally need to merge the two images into one layer by selecting “Flatten Image” under the “Layer” menu.

If the only image you have is an overexposed one, you can use your image pro-cessing program to mimic the effect of a coronagraph or solar eclipse, by blacking out the bright solar disc and showing just the prominences. Select the overexposed disc using “Magic Wand.” Then, under “Enhance,” then “Brightness/Contrast,” then “Levels,” move the right-hand pointer on the “Output Levels” as far to the left as possible. As you move the pointer to the left, the washed-out disc will darken until it is totally black, leaving you with a picture that looks as if the Sun has been blocked out by the Moon or the occulting disc in a coronagraph (Fig. 8.9 ).

Fig. 8.9 The webcam image of the Sun shown in Fig. 8.6 , with the overexposed solar disc blacked out using Adobe Photoshop Elements , mimicking the effect of a coronagraph or a total solar eclipse

197L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0, © Springer Science+Business Media New York 2012

Building a Solar Projection Box

As discussed in Chap. 2 , a projection box offers some distinct advantages for the solar observer. It allows the Sun’s image to be viewed safely without the inconvenience of having to hold a screen behind the telescope. It leaves your hands free to make notes, drawings and adjustments. It is much easier to make accurate records of sunspot posi-tions with a projected image than it is by looking directly at the Sun through a fi lter. If a projection box is attached to the telescope, the telescope’s fi eld of view does not move relative to the projection screen, and sunspots and other solar features can be observed and recorded very accurately. A projection box also allows the Sun to be observed by groups of people and, if thoughtfully constructed, can prevent curious children from putting their heads to the eyepiece (or at least make it very dif fi cult!).

The projection box described here was built entirely from 6.4 mm (1/4 in.) thick balsa wood bought from an art and crafts store. The interior of the box was lined with thin black card (also obtainable from art stores) in order to reduce internal re fl ections, which can spoil the contrast of the image. Two sides of the box are parallel, so that the pieces could be joined together easily using balsa wood framing of square cross-section, 12.5 mm (1/2 in.) on a side. However, the sides of the box were made so that they taper down to 75 mm (3 in.) at the eyepiece end, as other-wise the front of the box would collide with parts of the telescope tube and the whole thing would look bulky and ungainly. Note that because balsa wood is intended for making such things as model airplanes and other small objects and not large pieces of apparatus like projection boxes, it often comes in strips no more than 75 mm (3 in.) or 100 mm (4 in.) wide. If this is the case with your balsa wood, you will need to cut two or more pieces of wood and then glue them together. Glue all the joints using ordinary wood adhesive from a local hardware store.

The most important dimension to establish when building a projection box is its length. Using an eyepiece of a given magni fi cation, moving the screen further

Appendix A

198 Appendix A

behind the eyepiece increases the size of the projected image. To do serious solar work you need an eyepiece that will show the whole solar disc comfortably within the fi eld of view. Avoid using an eyepiece that barely accommodates the whole Sun and so shows the limb of the Sun’s disc very close to the edge of the fi eld, as the view at the edge of the fi eld is distorted in most eyepieces. On the other hand, using too low a power gives a very long projection distance for a good-sized disc and results in an inconveniently long projection box. An eyepiece giving a magni fi cation of between 50× and 65× gives the best results. The standard diameter used by ama-teur astronomers for a projected solar image is 152 mm (6 in.), although if you use a very small telescope you may fi nd a 100 mm (4 in.) disc easier to work with.

To fi nd the correct length for a projection box on your own telescope, draw a circle of your chosen disc size on a sheet of plain paper. Here we will use a 152 mm (6 in.) disc size as an example. Insert your choice of projection eyepiece in the tele-scope so that it is not quite fully placed in the drawtube but the eyepiece barrel sticks about 3 mm out. Determining the correct projection distance for a 152 mm disc should in theory be a simple matter: just measure the distance behind the eyepiece you have to hold the paper in order for the Sun’s image to exactly fi ll the circle.

You should note, however, that Earth’s orbit around the Sun is slightly elliptical, which means that Earth’s distance from the Sun varies during the course of the year. Earth is furthest from the Sun in July and closest in early January. 1 For the solar observer, the upshot of this is that the Sun’s apparent diameter is slightly larger in January than it is in July. More precisely, its January diameter is 32.5 arc min, whereas in July it is just 31.5 arc min. The average between these two extremes is 32 arc min, so the variation of 0.5 arc min either way means that the diameter of the projected image varies by about 1.56 % either way, which for a 152 mm disc translates into a variation of about 2.4 mm on either side. This is enough to make the real solar image too small for a 6-in. circle in summer and too large in winter, and will compromise the accuracy of solar drawings. Therefore, you need to allow the projection distance to vary by a corresponding percentage.

The easiest way to determine the maximum and minimum projection distances needed throughout the year is to simulate the variation in the Sun’s apparent diam-eter by drawing three concentric circles: one of your chosen disc diameter (in our case 152 mm), one of the maximum disc diameter (152 mm + 3 mm = 155 mm) and a third of the minimum diameter (152 mm − 3 mm = 149 mm). The increment has been rounded up from 2.4 to 3 mm in order to give slightly more tolerance than you actually need. You can then simply measure, and note down, the projection distance required to fi ll each circle. It is essential to measure the projection distance accu-rately: it is de fi ned as the distance between the rear surface of the eyepiece and the surface of the projection screen . In the case of the given setup, with an 80 mm

1 This may seem paradoxical to inhabitants of the northern hemisphere, who experience winter when Earth is closest to the Sun! But the cooling of the northern hemisphere in winter caused by the tilt of Earth’s axis away from the Sun totally overwhelms the very slight heating effect caused by Earth’s closer proximity to the Sun at this time of year.

199Appendix A

refractor of focal length 910 mm projecting using a 15 mm Plössl eyepiece (magni fi cation 61×), the distances were:

Average 263 mm Maximum 270 mm Minimum 258 mm

You need to build your box so that it allows for the maximum projection distance required. If the box is a little too long you can shorten the distance slightly, and thus get the image down to the correct size by pulling the eyepiece out a little in its drawtube, but if it is too short you cannot extend the box! The fi nal length for your projection box is this distance (marked d in Fig. A.1 ) plus the distance by which the

6.4 mm–thickBalsaWood

CornerJoint

ProjectionScreen

Eyepiece

d

Pipe Bracket

Star Diagonal(seen from behind)

Fig. A.1 Diagram showing the construction of the author’s projection box. The projection distance d determines the diameter of the projected solar image

200 Appendix A

eyepiece protrudes into the box. When you use the box, you can make small changes to the projection distance to allow for the variation in the Sun’s apparent size by adjusting the position of the eyepiece in its drawtube. Inserting the eyepiece fully will give the maximum projection distance, whereas pulling it out by a few millimeters and then clamping it again will give the minimum.

The screen end of the box can be square. When determining its size, allow for 5 mm of clearance around the solar disc (i.e., 5 mm on each side if measuring its cross section), plus the thickness of the square corner joints holding the box together (another 26 mm or 1 in. in total) and the thickness of the sides of the box (total 13 mm or ½ in.). For the box in the example, then, the total width of the base comes to:

152 mm 5 mm 5 mm 26 mm 13 mm 201 mm+ + + + =

Thus the screen end of the box should be 201 mm square. The top piece, through the center of which the eyepiece mount should be inserted, has the same width as the tops of the side pieces (75 mm, 3 in.) and the same length as the internal width of the box (188 mm). The hole for the eyepiece or drawtube (see below) should be just wide enough for the tube to slide through with as little play as possible. If you have access to a lathe or other precision cutting device, boring a hole of the correct diameter should not be a problem. If you only have access to hand tools, you can make a very functional eyepiece hole quite easily by drawing a circle of the correct diameter in the exact center of the top piece and then drilling a series of small holes around its circumference. If the holes are spaced reasonably closely together, you can then break down the walls between them with a chisel until the central piece of wood falls out. You can fi le the interior of the resulting hole smooth with sandpaper—though be careful not to make it too big for the tube. The front of the box—i.e., the side not used for looking at the Sun’s image—needs to be blocked in, preferably with more balsa wood for rigidity, or with thick black card. You now have a box with just one open side, into which you look to view the Sun’s projected image.

How you attach the box to the telescope depends on the characteristics of your instrument. Ideally, the box should slide over the drawtube and be clamped either with a screw attached to a collar at the front of the box, or by making the box fi t very tightly over the drawtube (e.g., by lining the eyepiece hole with felt). However, some telescopes, particularly small refractors, may not reach focus without a star diagonal. In this case, you need to slide the box over the eyepiece end of the star diagonal, so that the box’s long dimension is at right angles to the telescope tube. Clamping the box rigidly to the telescope tube requires some ingenuity if you use a star diagonal. I attached a pipe bracket (obtained from the plumbing section of a hardware store) to the top piece of the box, with a bolt passing through it and so clamping the box tightly to the eyepiece collar of the main telescope. Two smaller bolts going through the sides of the pipe bracket keep the box’s orientation steady (Figs. A.1 and A.2 ).

201Appendix A

Using the star diagonal method of projecting the Sun’s image is probably better than direct projection, as you can better shade the interior of the box from ambient sunlight, and the near-horizontal position of the Sun’s image makes it easier to view. A disadvantage is that it projects a reversed solar image—i.e., east left, west right, which is the opposite way around from a solar image projected straight through. Orientation of the Sun’s image is discussed in more detail in Chap. 4 .

To obtain the best projected images you should use a good quality, thick paper. Many amateurs over the years have used “Bristol board,” a very smooth white paper available from art shops. This gives very good results, but the thick, smooth car-tridge paper with a slightly creamy tint shows sunspot detail even better. Faculae, though, are more prominent on Bristol board. You can use a sheet of Bristol board and a sheet of cartridge paper pasted together, so that you can fl ip between the two surfaces to see the two types of solar features to their best advantage. To plot sun-spot positions use a rotatable projection grid made of Bristol board; again, this is discussed in detail in Chap. 4 .

Fig. A.2 The author’s projection box, attached to an 80 mm (3.1 in.) refractor with a pipe bracket

203L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0, © Springer Science+Business Media New York 2012

Equipment Suppliers

This list is not intended to be comprehensive but rather is a selection of suppliers of solar observing equipment that amateur astronomers have found to be useful. It does not constitute a recommendation of these suppliers and neither the author nor the publisher makes any guarantee as to the safety of these products or their suit-ability for solar observing. Try to solicit the advice of an experienced solar observer before buying any piece of equipment.

Suppliers of astronomical equipment exist in vast numbers. This list is restricted, therefore, to those companies selling products either speci fi cally for solar observing or useful for solar work. To fi nd suppliers of telescopes and other general astro-nomical equipment, try your favorite Internet search engine or peruse the advertise-ments in astronomy magazines (see Appendix D , “Further Reading”).

Suppliers in North America

Astro-Physics, Inc. , www.astro-physics.com . Suppliers in USA of Baader Planetarium fi lters and accessories.

Celestron International , www.celestron.com . Solar fi lters and accessories to suit the wide range of Celestron telescopes; also the NexImage camera, a webcam optimised for astronomy. Celestron do not sell directly to the public, but their prod-ucts can be obtained through Celestron dealers worldwide. Contact Celestron to fi nd your nearest dealer.

Coronado , part of Meade Instruments, www.meade.com/product_pages/coro-nado/coronado.php . A leading manufacturer of sub-angstrom H-alpha fi lters and

Appendix B

204 Appendix B

solar telescopes for the amateur solar observer, including the Personal Solar Telescope (PST), a complete H-alpha solar telescope for just under $500.

DayStar Filters LLC , www.daystar fi lters.com . The longest-running manufac-turers of sub-angstrom H-alpha fi lters.

Kendrick Astro Instruments , www.kendrickastro.com/astro/solar fi lters.html . Suppliers of Baader AstroSolar fi lters.

Lumenera Corporation , www.lumenera.com . Top-of-the-line webcam-type cameras suitable for high-resolution solar imaging.

Orion Telescopes & Binoculars , www.telescope.com . Suppliers of full-aper-ture glass solar fi lters for white-light observing and Coronado H-alpha telescopes; also many useful accessories, including adapters for mounting digital cameras to telescopes.

Point Grey Research , www.ptgrey.com/products/ fl ea3/ fl ea3_ fi rewire_camera.asp . Suppliers of the tiny, but highly effective, Flea3 webcam-type camera suitable for high-resolution solar imaging.

Thousand Oaks Optical , www.thousandoaksoptical.com . Manufacturers of both glass and plastic white-light solar fi lters for visual observing, solar eclipse viewers, glass fi lters for solar photography and 1.5-Å H-alpha prominence fi lters.

Suppliers in the UK and Europe

Telescope House (Broadhurst Clarkson and Fuller Ltd), www.telescopehouse.com . Suppliers of white-light solar fi lters and Coronado H-alpha fi lters and telescopes.

Baader Planetarium , www.baader-planetarium.com . Manufacturers of Baader AstroSolar Safety Film and AstroSolar Photo Film Mylar-type aperture fi lters, also many other fi lters and accessories for solar observing—e.g. the “Fringe-Killer” fi lter.

David Hinds Ltd , www.dhinds.co.uk . UK suppliers of Baader Planetarium equipment; also of telescopes and accessories by Celestron. See the two separate websites: www.celestron.uk.com and www.baader-planetarium.uk.com .

The Imaging Source , www.theimagingsource.com/en_US . Manufacturers of advanced webcam-type cameras for high-resolution imaging.

Lunt Solar Systems , www.lunt-solarsystems.eu . Manufacturers of a wide range of H-alpha and Ca-K solar telescopes and fi lter systems.

SCS Astro , www.scsastro.co.uk . Solar fi lters, including DayStar fi lters. Solarscope , www.solarscope.co.uk . Manufacturers of high-end H-alpha solar

fi lters and telescopes. The Widescreen Centre , www.widescreen-centre.co.uk . UK suppliers of Lunt

solar fi lters and telescopes.

205L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0, © Springer Science+Business Media New York 2012

Solar Observing Organizations

Astronomy Clubs

If you are new to solar observing, and certainly if you are starting out in astronomy, the fi rst place you should go for advice on observing is your nearest local astron-omy club or society. There are bound to be some members there with experience of solar observing and most are willing to provide help and encouragement as well as advice on what equipment to buy. If you cannot fi nd a local society in your area, try contacting the following national organizations of astronomy clubs:

(In the USA): The Astronomical League, www.astroleague.org/ , maintains a list of local astronomy clubs and their contact details. The monthly magazine Sky & Telescope (see also under “Magazines” in Appendix D ) has a web page that allows you to search for astronomy clubs (and also museums, observatories and planetari-ums) in your area: http://www.skyandtelescope.com/community/organizations .

(In the UK): The Federation of Astronomical Societies (FAS) , http://www.fedastro.org.uk/fas/ , maintains a list of local societies that are members of the FAS, including their websites and details of their meetings.

National Organizations in the United States

The Solar Section of the American Association of Variable Star Observers , www.aavso.org/solar , focuses on determining the American Relative Sunspot Number ( R

A ) using sunspot counts supplied by contributing members.

Appendix C

206 Appendix C

The Association of Lunar and Planetary Observers , http://alpo-astronomy.org/ , also has a Solar Section which emphasizes recording solar activity pictori-ally—by drawings and electronic imaging.

Royal Astronomical Society of Canada , www.rasc.ca . Publishes the very use-ful annual Observer’s Handbook .

National Organizations in the UK

The leading association of amateur astronomers in Great Britain is the British Astronomical Association , www.britastro.org . The BAA has a very active Solar Section, to which several dozen members send monthly observations. You can access its website through the BAA home page above.

If you are new to solar observing you may wish to consider joining the Society for Popular Astronomy , www.popastro.com . The SPA is a national society for beginning and intermediate amateur astronomers of all ages. It has a very active Solar Section that helps beginners master the basic techniques of solar observing. The SPA also publishes a lively bimonthly magazine, Popular Astronomy ; in this, the Solar Section publishes regular reports based on observations sent in by members.

207L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0, © Springer Science+Business Media New York 2012

Further Reading

Books on the Sun and Solar Observing

BECK, R. et al., Solar Astronomy Handbook (Willmann-Bell, 1995). A very detailed compendium of amateur methods of observing and imaging the Sun, cov-ering a number of technical topics not described in this book, including observing the Sun at radio wavelengths. Rather dated now, but a good reference for the more advanced solar observer.

GOLUB, L. and PASACHOFF, J. M., Nearest Star: the surprising science of our sun (Harvard University Press, 2001). An easy-to-read guide to how the Sun works.

JENKINS, J. L., The Sun and How to Observe It (Springer-Verlag, 2009). A technical guide to observing and imaging the Sun.

LANG, K. R., Sun, Earth and Sky (2nd edition, Springer-Verlag, 2006). A guide to our current knowledge of the Sun and its in fl uence on Earth.

LANG, K. R., The Cambridge Encyclopaedia of the Sun (Cambridge University Press, 2001). A beautifully-illustrated guide to our nearest star. It describes the workings of the Sun to a quite high technical level, although mathematics are pre-sented separately in text boxes and so do not interrupt the fl ow of the narrative.

MOBBERLEY, M., Lunar and Planetary Webcam User’s Guide (Springer-Verlag, 2006). An introduction to webcam imaging; includes a chapter on imaging the Sun with a webcam.

PHILLIPS, K. J. H., Guide to the Sun (Cambridge University Press, 1992). An accessible but detailed guide to the Sun and how it works.

Appendix D

208 Appendix D

PUGH, P., Observing the Sun with Coronado TM Telescopes (Springer-Verlag, 2007). A survey of the various solar telescopes, fi lters and accessories that were available at the time of the book’s publication.

TAYLOR, P. O., Observing the Sun (Cambridge University Press, 1991). A detailed guide to monitoring solar activity by sunspot counting and electronic methods.

Books on Photography and Digital Imaging

BERRY, R., and BURNELL, J., The Handbook of Astronomical Image Processing , 2nd Edition (Willmann-Bell, Inc., 2004). A detailed guide to image processing, with a companion CD-ROM containing AIP4WIN image processing software.

COVINGTON, M. A., Astrophotography for the Amateur (2nd Edition, Cambridge University Press, 1999). A detailed guide to all aspects of amateur astrophotography.

COVINGTON, M. A., Digital SLR Astrophotography (Cambridge University Press, 2007). A specialised guide to astrophotography techniques with DSLR cam-eras. Contains some very brief notes on webcam imaging.

DRAGESCO, J., High Resolution Astrophotography (translated by Richard McKim, Cambridge University Press, 1995). A “classic” from the 35 mm fi lm era, this book is nevertheless essential reading for the serious solar observer and imager today, for it includes thorough technical discussions on resolution, seeing condi-tions, observing sites, telescopes and accessories—all as important in the digital age as they were in the age of fi lm.

IRELAND, R. S., Photoshop Astronomy (2nd edition, Willmann-Bell, Inc.). A detailed guide to astronomical image processing using Adobe Photoshop . Comes with a companion DVD.

REEVES, R., Introduction to Digital Astrophotography: Imaging the Universe with a Digital Camera (Willmann-Bell, Inc., 2005). A superb guide to imaging with digital cameras. It covers compact digital cameras as well as DSLRs and has a strong chapter on webcam imaging.

REEVES, R., Introduction to Webcam Astrophotography: Imaging the Universe with the amazing, affordable webcam (Willmann-Bell, Inc.). A book devoted to webcam imaging, by the author of Introduction to Digital Astrophotography (above).

Reference Books

The Handbook of the British Astronomical Association is published annually by the British Astronomical Association (Burlington House, Piccadilly, London). It con-tains detailed solar data, including P , B

0 and L

0 (essential for working out solar

co-ordinates) tabulated at 5-day intervals. See Appendix B for more about the British Astronomical Association.

209Appendix D

The Astronomical Almanac is also published annually and is a collaboration between Her Majesty’s Stationery Of fi ce in the UK and the United States Naval Observatory. It is available in both countries from astronomy book suppliers. The Almanac tabulates P , B

0 and L

0 at daily intervals. Willmann-Bell, Inc. also pub-

lishes a software version of the Almanac known as the Multiyear Interactive Computer Almanac 1800–2050 .

Magazines

The major commercial astronomy magazines often contain articles on solar observ-ing and the latest developments in our understanding of the Sun. If you are a serious solar observer (or a serious amateur astronomer of any kind), it is a good idea to get at least one magazine regularly, as they often present information that is more up-to-date than in books.

Sky and Telescope ( www.skyandtelescope.com/ ) contains many “how to” observing articles. Astronomy ( www.astronomy.com/ ) also has many articles on practical astronomy. Both magazines can often be found in British newsagents as well.

In the United Kingdom, Astronomy Now ( www.astronomynow.com/ ) and Sky At Night ( www.skyatnightmagazine.com/ ), published by the BBC, are both avail-able from many British newsagents. Both are very strong on equipment reviews.

Websites

The following websites may be of use to the solar observer:

Current Solar Images From Earth Big Bear Solar Observatory: http://www.bbso.njit.edu/ Mount Wilson Observatory: http://www.mtwilson.edu/sci.php Uccle Solar Equatorial Table (images from Europe): http://sidc.oma.be/uset/

index.php

Current Solar Images From Space SOHO spacecraft: http://sohowww.nascom.nasa.gov/ Solar Dynamics Observatory: http://sdo.gsfc.nasa.gov/ STEREO spacecraft: http://stereo.gsfc.nasa.gov/

Solar Activity, Space Weather and Aurora Warnings Current solar activity and space weather: http://spaceweather.com/ (also includes

other astronomical images sent in by viewers). Space Weather Prediction Center: http://www.swpc.noaa.gov/ Current solar active regions with of fi cial AR numbers: http://www.nwra.com/

spawx/listsrs.html

210 Appendix D

Solar Position Measurements Peter Meadows’s website (UK)—Stonyhurst discs and software for determining

heliographic coordinates: www.petermeadows.com

TiltingSun —software for determining P , B 0 and L

0 and showing the Sun’s current

tilt and orientation: http://www.atoptics.co.uk/tiltsun.htm Current values of P , B

0 and L

0 : http://www.jgiesen.de/sunrot/index.html

Sunspot Numbers Solar In fl uences Data Analysis Center (SIDC)— http://sidc.oma.be/ . Publishes

the of fi cial Relative Sunspot Number, in continuation of the Relative Sunspot Number begun by Rudolf Wolf in 1848.

American Association of Variable Star Observers (AAVSO)—Solar Section— http://www.aavso.org/solar . Issues a monthly electronic Solar Bulletin that publishes the American Sunspot Number— R

A .

211

A Active areas , 84–88, 90, 91, 93, 94, 103, 128 Active prominences , 125–127 Active regions , 9–11, 55, 59, 79, 105, 115,

118, 125–127, 130, 131, 182 Adapters (for cameras) , 143, 150, 154, 168 Adobe Photoshop , 187, 193, 194, 196 Adobe Photoshop Elements , 187–189,

193–196 Afocal method (photography) , 151, 164, 168,

169, 190 Alt-azimuth mountings , 25, 26, 69, 135 Aperture fi lters , 24, 25, 32, 33, 35–37, 45, 98,

146, 152, 156, 163, 168, 180 Aurorae , 13, 14, 55, 61 Autofocus , 163, 166 Automatic exposure , 141, 162

B B

0 , 73–76, 90, 97, 98

Baader coronagraph , 119–121 Baader fi lters (white-light) , 34, 125 Baader K-line fi lter , 125 Binoculars , 37–39 Bipolar sunspot group , 51, 53, 75 Brightness (in image processing) , 88, 89,

191–194, 196 Butter fl y diagram , 78, 79

C Cable release , 144, 155, 157, 164 Ca-K. See Calcium-K Calcium-K , 62, 102, 124–125 Camera adapters , 148, 150, 154, 157, 158, 168 Cameras

CCD , 63, 123, 137, 138, 145, 169, 177, 178, 194

digital compact , 137–144, 146, 148, 151, 162–165, 168–170, 173, 175, 176, 178, 184, 186, 188, 192

digital SLR , 103, 141–144, 164 webcam-type , 21, 169, 176, 178, 198

Carrington, R. , 61, 70 Catadioptric telescopes , 22, 24–27,

35–37, 145 Central meridian , 74, 76–78 Chromosphere , 5, 6, 9, 10, 22, 101–134,

168–172, 181 CMEs. See Coronal mass ejections (CMEs) Colour (of digital image) , 194 Compact cameras. See Cameras Compact fl ash cards. See Memory cards Contrast (in image processing) , 33–35,

38, 44, 47, 48, 52, 97, 98, 108, 111, 112, 115, 117–120, 144, 145, 150, 151, 160, 164, 168, 169, 171, 183, 185, 186, 188–194, 196

Index

L. Macdonald, How to Observe the Sun Safely, Patrick Moore’s Practical Astronomy Series,DOI 10.1007/978-1-4614-3825-0, © Springer Science+Business Media New York 2012

212 Index

Convective zone (of sun) , 4, 47 Core (of sun) , 3, 43 Corona , 5, 6, 10–12, 16, 101, 103, 105, 107,

119, 127, 153 Coronado CEMAX eyepieces , 116 Coronado fi lters , 111, 115 Coronagraph , 6, 101, 103, 107, 108, 119–121,

171, 172, 196 Coronal holes , 11, 13, 14 Coronal mass ejections (CMEs) , 11–13,

15, 16 Cropping , 187, 189, 190

D DayStar fi lters , 108, 109, 113, 114, 116–119,

124, 125 Deslandres, H. , 107 Dobsonian mounting , 25, 151 DSLR cameras. See Cameras

E Eclipse (of moon) , 5, 6, 119, 120, 153, 167, 174 Eclipse (of sun) , 4–6, 10, 11, 32, 34, 37, 38,

101, 107, 119, 120, 153, 174, 196 Energy rejection fi lter (ERF) , 109, 116, 117,

119, 124 Equatorial mounting , 20, 25, 26, 68–70, 80,

120, 145, 179 ERF. See Energy rejection fi lter (ERF) Eruptive prominences , 126 Etalon , 108, 109, 111–119, 122, 124 Exposure (in solar photography) , 24, 135, 137,

141, 145, 146, 150, 159, 163 Eye, danger to from Sun , 17, 18 Eyepiece projection , 148, 150, 156–160, 168,

173, 180 Eyepieces (for solar projection) , 30, 31, 89

F Faculae, polar. See Polar faculae Filaments , 50, 51, 104, 105, 107, 111, 115,

118, 119, 125–131, 169, 170 Film, photographic , 32, 63, 138 Filters

Baader (white light) , 34, 125 Ca-K , 106, 124, 125 eyepiece (DANGER!) , 18, 22, 31, 36 glass , 25, 34, 35, 37, 60, 98, 145, 147 H-alpha ( see H-alpha fi lters) Mylar , 32–35, 39, 59, 98, 122, 145, 168, 193

photographic fi lm (DANGER!) , 32, 35, 147, 148, 168

smoked glass (DANGER!) , 32 in solar photography , 24, 34, 135, 137,

145–147 Flares, solar , 9, 15, 61, 105, 126, 131 Flocculi , 105 Focusing (for photography) , 123, 135, 142,

143, 147, 160, 164–166, 169 Focusing magni fi ers , 142, 166 Focusing screen (in DSLR cameras) , 164, 165,

167, 170, 175 Fork mountings , 26 Fringe-Killer fi lter , 31, 168, 180, 194

G Geomagnetic storms , 13–16, 134 Glass fi lters , 34, 35, 60, 145, 193 Go To telescopes , 22, 26 Granulation, solar , 31, 47 Graphs of solar activity , 94

H Hale, G.E. , 107 H-alpha. See Hydrogen-alpha H-alpha fi lters , 22, 93, 105–110, 113–126,

128, 130, 131, 133, 135, 137, 169, 181, 183

Hedgerow prominences , 126 Herschel wedge , 36, 37, 44, 125, 144, 145,

168, 179 Hodgson, R. , 61, 62 Hydrogen-alpha (H-alpha) , 5, 60, 80, 93, 102,

135, 174

I Image processing software , 174, 177, 186, 187 Imaging Source (webcam-type cameras) ,

177, 178 Interference fi lters , 47, 102, 107–110, 168, 180

J Janssen, P. , 107

L L

0 , 74, 76–78

Light bridge , 54, 62, 88, 93, 98 Limb darkening , 48–49, 53, 59, 140, 159, 191

213Index

Little Ice Age , 13 Lockyer, N. , 107 Lumenera (webcam-type cameras) , 177, 178 Lunt fi lters , 116, 122 Lyot, B. , 107

M Maksutov telescope , 24, 25, 27, 31, 35, 68, 69,

84, 144, 153, 154 Maunder, E.W. , 13, 78 Maunder minimum , 13 McIntosh sunspot classi fi cation system , 57 MDF. See Mean Daily Frequency Mean Daily Frequency , 84–86 Memory cards , 139, 160, 173 Motor drives (for telescopes) , 28, 67, 135,

166, 167 Mylar fi lters , 32–35, 193

N Naked-eye sunspots , 37, 98–99 Newtonian telescope , 23, 157 NexImage (webcam by Celestron) , 176, 177,

180, 182, 183

O Off-axis mask , 23–24 Orientation of sun’s image , 26, 69, 70, 72,

129, 191

P P . See Position angle Passband , 108–119, 124, 125, 168 Penumbra , 50–57, 62, 68, 86, 88, 89, 160 Penumbral fi laments , 50, 51 Photosphere , 4–7, 11, 48–51, 53, 54, 59, 62,

88, 101, 102, 169, 170, 191, 192 Plages , 104–106, 115, 118, 119, 131 Polar faculae , 59–61, 78, 97, 98 Pores , 53, 54, 56, 78, 86, 88, 89, 183 Position angle , 72–76, 90, 97 Pre- fi lter. See Energy rejection fi lter Prime focus , 30, 122, 123, 150–155, 160,

168, 180 Projection box(es) , 20, 28–31, 44, 45, 47, 60,

61, 64–67, 70, 86, 97 Projection grid , 64–67, 90, 97 Projection method , 23, 30, 31, 45, 61, 64–80,

89, 98, 128, 150, 190

Prominences , 10, 11, 22, 93, 102–105, 107, 108, 110, 111, 115, 117, 119–121, 125–131, 135, 168, 169, 171, 172, 180, 181, 183–186, 194–196

Q Quiescent prominences , 125–127, 130

R R . See Relative Sunspot Number Radiative zone , 3 Re fl ector (telescope) , 19, 22–25, 27, 31, 37,

68, 84, 151, 153, 180 Refractor (telescope) , 8, 19–31, 33, 37, 44, 46,

47, 50, 61, 68–70, 84, 87, 97, 103, 109, 112, 114, 115, 119, 120, 123–125, 135, 144, 145, 147, 153–159, 162, 165–171, 175, 176, 178–180, 183, 188–190, 194

Relative Sunspot Number , 83, 86–96, 103 Report forms , 91, 92, 94 Rotation number , 70, 74

S Satellites, effects of solar activity on ,

9, 11, 13, 15 Scanning images , 70 Schmidt-Cassegrain telescope , 24, 27, 31, 68,

69, 84, 117, 144 SCT. See Schmidt-Cassegrain telescope Seeing , 5, 14, 42–44, 48, 50, 61, 62, 70, 80,

87–89, 93, 97–99, 106, 120, 122, 136, 138, 145, 146, 159, 160, 173–175, 179, 180, 182–184, 190

Self-timer , 141, 164 Sharpening digital images , 192–193 Slow motions (on telescope mount) , 20, 26,

47, 67, 68, 115, 145, 166 Solar and Heliospheric Observatory (SOHO) ,

9, 11, 12 Solar cycle , 8, 9, 11, 14, 16, 50, 58, 78, 79, 84,

90, 94, 96, 97, 102, 131 Solar Dynamics Observatory (SDO) , 9 Solar fi nderscopes , 45, 46 Solarscope fi lters , 109, 113, 114, 116, 122, 123 Solar telescope , 34, 106, 108, 110–112, 115 Solar wind , 6, 11–13, 15 Spectroheliograph , 107 Spectrohelioscope , 107, 108 Spectroscope , 107 Spicules , 127, 181

214 Index

Spörer, G. , 78 Spörer’s law , 78, 90 Stars (compared with sun) , 1–3, 6, 17,

22, 120, 171 STEREO spacecraft mission , 16, 55 Stonyhurst discs , 71, 72, 74, 76, 77, 86, 90 Sun diagonal. See Herschel wedge Sunspot cycle , 7–9, 13, 70, 71, 78, 79, 96,

97, 125 Sunspot groups , 7, 9–11, 15, 37, 47–51,

53–62, 75, 76, 78–80, 83–87, 89, 93, 96, 98, 99, 104, 105, 118, 126–128, 130–135, 145, 146, 148, 156, 158, 159, 166, 174, 186

bipolar , 51, 53, 55–57, 75 Sunspots

classi fi cation , 55–59 longevity , 54 magnetic fi elds , 59 numbers , 53, 55, 59 positions , 63, 64, 67, 71–74 temperature , 7, 13

T T-adapter , 154 Teleconverter , 152, 154–156, 159, 160, 167,

169, 188, 189

Tele-extender , 158 Transition region , 6, 9 Transition Region and Corona

Explorer(TRACE) spacecraft , 9 T-ring , 120, 154, 156–158, 164

U Umbra , 50–56, 77, 86, 88, 93, 99 Unsharp mask , 189, 193

W Weather, sun’s effect on , 1, 12, 13,

15, 43, 44 Welder’s glass , 37, 99 White-light fl ares , 61, 62, 80, 93,

97–98, 131 Wilson, A. , 52 Wilson Effect , 52, 53 Wolf, J.R. , 87 Wolf Number. See Relative Sunspot Number

Z Zurich Number. See Relative Sunspot

Number