instrumentation concepts ground-based optical telescopes

Instrumentation Concepts

Ground-based Optical Telescopes

Keith Taylor(IAG/USP)

Aug-Nov, 2008

Aug-Sep, 2008 IAG-USP (Keith Taylor)

Imaging considerations

Trading field of view vs. angular resolution A large field at coarse spatial resolution or

smaller field of view at high fidelity? Generally detector defines fixed pixel format

Constraints also driven by seeing and pixel scale (or camera f-ratio)

Photometric accuracy Simple morphological discrimination or accurate

flux measures as well? Definition of passbands and central

wavelengths Multiple simultaneous passbands?

Astronomical CCD Imaging

Simplest astronomical instrument (in principle)

Undispersed 2D images of field of view

Generally use filters to limit spectral band-pass Detector itself may

supply band-width Polarimetric capabilities?

Rapid reads can give limited time resolution information

Standardized Filter Systems

Variety of different filter systems prevalent in optical/IR domain. eg: UBVRI / JHKLM – Johnson/Cousins

(UV/optical) (NIR) u g r i z - Sloan Digital Sky Survey (SDSS)

optical bands Extras and modifications

y - UKIDSS IR band K’ and K* - modifications to K to avoid thermal

radiation HST, Spitzer etc. defined by wavelength

rather than name (revolutionary!)

Examples of CCD Imaging

CCD = Charge-coupled deviceMore sensitive than photographic plates by factors ~50

Data can be read directly into computer memory, allowing digital enhancement and manipulation

Negative image can enhance contrasts

False-color image to visualize brightness contours

cf: Photographic Imaging(eg: AAT c1990)

14inch (=350mm)• ~5 Giga Pixels

cf: Largest CCD currently:• 8 Mega Pixels

(4k-by-2k)• CCD arrays up to

~0.1 Giga Pixels

Generally seeing-limited

• How do we obtain higher spatial definition?

Adaptive OpticsComputer-controlled “adaptive” mirror adjusts the mirror surface (many times per second) to compensate for distortions induced by atmospheric turbulence

And yet furtherspatial resolution?

InterferometryRecall: Resolving power of a telescope depends on diameter D:

min = 1.22 /D.

This holds true even if not the entire surface is filled out.

Sparsely filled aperture:

•Combine the signals from several smaller telescopes to simulate one big mirror

Interferometry

Spectroscopic considerations

What kind of spectral feature are of interest? Emission or absorption lines; continuum

shapes Broad, narrow or spectrally unresolved Low or high contrast with continuum Spectral Energy Distributions (SEDs)

Line centres ; equivalent widths ; line shapes ; kinematic mapping? and/or precise spectrophotometry?

One or many targets simultaneously?

The simplest spectrographUsing a prism (or a grating), light can

be split up into different wavelengths (colors!) to produce a spectrum.

Spectral linesSpectral lines in a spectrum in a spectrum tell us about the chemical tell us about the chemical composition and other composition and other properties of the observed properties of the observed object object

Typical grating spectrograph

Simple grating spectrograph

Spectrum extracted along a slit so ‘imaging’ in one dimension

Off source light along slit used to measure and subtract sky background

What you get Optical long slit

spectrum of a galaxy Minimal data

reduction in displayed spectral image

Can see galaxy, bad pixels, cosmic ray hits and sky lines Need off source

signal to measure and extract target (sky subtraction)

Sky lines

Target

Considerations for Spectroscopy

Basic parameters - resolution and central wavelength for spectrum

Slit width (if selectable) affects resolution Wavelength range

Set by combination of detector geometry and spectral resolution

Some spectrographs provide large -range at low-R; others provide only a few 1,000kms-1 range, so centering on a critical line of interest (eg: H)

But, what if you need both high-R and large -range?

High Resolution and lots ofSpectrum

• X-dispersed echelle grating spectrometers allow high resolution and lots of spectral coverage• Achieve this by having two

orthogonal gratings• One gives the high resolution (in y-axis) the other spreads the spectrum across the detector(in x-axis)• However, the slit is consequently much shorter

STELES echelle spectrograph(for SOAR)

Primary disperser (echelle grating)

Secondary (orthogonal) disperser (VPHG)

Redchannel

Bluechannel

Multiobject Spectroscopy

To get spectra for lots of objects at once. Can use two approaches Multislit - have several slits in the

image plane and get spectra for all of them

Use fibres to pipe light from different parts of the focal plane while reformatting them along the spectrograph slit

Both techniques were developed in the 80s and perfected in the 90s

Fibre Fed Systems

AAT 2dF (now AAT 2dF (now replaced by replaced by AAOmega)AAOmega) Pickoff fibres Pickoff fibres

positioned by robotpositioned by robot Include sky fibres Include sky fibres

for each objectfor each object

Multi-slit spectroscopy

Example of multislit spectral image

Easier to achieve at telescope (can use holes in a mask) but preparation and reduction can be more complex

Need to ensure spectra don’t overlap

LDSS-2mask

superimposed on sky image

Great care has to be taken in selecting objects to study so that they don’t overlap in wavelength direction.

Also need objects of similar brightness so the SNRs are similar.

Mask optimization is NOT trivial!

Field acquisition is NOT trivial

But what if you want images and spectroscopy simultaneuously?

Integral Field Spectroscopy Extended (diffuse) object with lots of spectra

Use “contiguous 2D array of fibres or ‘mirror slicer’ to obtain a spectrum at each point in an image

SIFS

Tiger

MPI’s 3D

Large-field imagingspectrographs

Narrow band filters Image a field in a single narrow band Use enough narrow bands and you

have very low res. spectroscopy Fabry-Perot

Effectively acts as a narrow tunable filter

Can thus image a field in emission lines of choice (eg. TTF)

Fabry-Perot Light enters etalon and is

subjected to multiple reflections

Transmission spectrum has numerous narrow peaks at wavelengths where path difference results in constructive interference need ‘blocking filters’ to use as

narrow band filter Width and depth of peaks

depends on reflectivity of etalon surfaces: finesse

Fourier Transform Spectrometer

• As translation mirror scans an interference pattern is produced that is the FT of the source spectrum• Scan distance defines the resolution of the spectrum• Advantage - get spectrum of whole field• Disadvantage - get broad band noise

IFTS for NGST

Detectors for Opical/near-IR

(current) Photon Counters:

Image tube + TV camera + real-time discrimination (not solid state)

eg: IPCS - c70s to c80s CCDs now dominate - Hi QE but …

Integrate signal on detector – no time resolution Finite read-noise Finite read-time

EMCCDs – new generation of Photon Counters CCD-like QEs V. high frame-rates

DQE - the key to gooddetectors

Detector quantum efficiency - the fraction of incident photons detected - is the key measure for the effectiveness of a detector;

Traditional photographic plates, while large in size, have DQE of only about 10%

CCDs and similar semiconductor devices can have DQE as high as 90% (though wavelength dependent) Like having a telescope with 9 times the

collecting area

CCDs

CCDs combine photon detection with integration and multiplexing

Incident photons excite charge carriers which are stored and integrated in a capacitor

CCDs are also uniquely effective in transferring charge from 2D to 1D charge ‘clocked’ from pixel to pixel and

read out at fixed point ideal for multiplexing

CCD Array Camera

Semiconductor fabrication limits the size of a CCD detector

To get a large area need to mosaic detectors together

Subaru Mosaic CCD Camera

Near-IR Detectors

CCDs use Silicon as their substrate Valance to conduction bandgap in silicon is 1.1eV so

restricted to detecting photons with wavelength < 1 micron

Need different materials for infrared InSb for 1 to 5 micron, HgCdTe for 1 to 2.5 micron Detector elements bonded to Si CCD system to provide

multiplexing readout

IR Arrays vs. Optical

IR arrays are smaller, more expensive (by factor of ~10/pixel)

Readout has to be faster because of higher backgrounds

Use of different materials can push to longer wavelengths More difficult to work with, less helpful

characteristics, more expensive At longest wavelengths have to stress the

detector to produce lower energy band gaps

Lecture 1 (2 sessions) - Synopsis

Fundamental Principles Introduction Information Theory Seeing-limited observations Diffraction-limited observations Signal-to-noise estimates

Lecture 2

Fundamentals

Aug-Nov, 2008 IAG/USP (Keith Taylor)

UVOIR Astronomy

Definition: UVOIR = the "UV, Optical, Near-Infrared"

region of EM spectrum Shortest wavelength: 912 Å (or 91.2 nm) --- Lyman

edge of H I; interstellar medium is opaque for hundreds of Å below here

Longest wavelength: ~3µm (or 3000 nm) --- serious H2O absorption in Earth's atmosphere above here

Ground-based UVOIR: 0.3µm (or 300nm) < < 2.5µm (or 2,500nm)

UVOIR Astronomy

Uniqueness: Best developed instrumentation; Best understood astrophysically; Highest density of astrophysical information; Provides prime diagnostics on the two most important

physical tracers.

===> UVOIR observations/identifications are almost always prerequisites to a

thorough understanding of cosmic sources in other EM bands.

Proof

Stars Plasma (to 105K)

Observational Priorities

Astronomy driven by discoveries rather than theoretical insights Direction of field shaped by observations in about

3/4 of instances. Few important astronomical discoveries were

predicted; many were actually accidental

Accidental Discoveries

Uranus Expanding universe Pulsars Supermassive black holes/AGN‘s Large scale structure Dark matter in spiral galaxies X-ray emitting gas in clusters of galaxies Gamma ray bursts Extra-solar planets High redshift evolution of galaxies

HST contributions were actually hindered by theoretical prejudice. A deep pencil-beam survey was delayed by 5 years.

Counterexamples: theory-driven discoveries Neptune General relativistic distortion of space-

time near Sun 21 cm line of HI Helioseismology Cosmic microwave background

ConclusionsIs Observational Astronomy a Science?

(Build, don’t Think)

Technology drives Discovery

Key technology development for UVOIR astronomy: 17th century: telescopes 19th century: spectroscopy, photography, quality

lens making, large structure engineering 20th century: large mirror fabrication,

electronic detectors, digital computers, space astronomy

Since 1980: array detectors

Telescope size: determines ultimate

sensitivity Diameter doubling time ~45 years Largest telescopes now 8-10m diameter

Collecting area of 10-m is 4*106 that of the dark-adapted eye

In planning: 15-m to 40-m For a given technology, cost D2.6

Cost is roughly proportional to mass Even using new technologies, next

generation of large ground-based telescopes will cross the $1 billion threshold.

The Future?

NB: Number of ground-based telescopes is NOT inversely proportional to their size Almost as many 8m

telescopes as there are 4m telescopes

How many 30m telescopes are there going to be in the next 50 years? (at US$1B a pop)

Review of some Basics• c = ν = 3.1010 cm/s• E = hν (ergs)• F = L/4d2

• G = 6.67.10-8 (c.g.s)• c = 3.1010 cm.s-1

• k = 1.38.10-16

• h = 6.626.10-27

• mH ~ mproton = 1.67.10-24 grams

• me = 0.91.10-27 grams• eV = 1.602.10-12 ergs• Luminosity of Sun = 4.1033 ergs/sec• Mass of the Sun = 2.1033 grams

Flux measurements Signal-to-Noise Ratio

"Sensitivity"---i.e. the faintest source measurable---is not simply a matter of the size of the photon collector.

It is instead a signal-to-noise (SNR) issue: SNR = measured value / uncertainty and is

dependant on many things, including: Structure of source (point vs. extended) Nature of luminous background & surroundings Foreground absorption Telescope & instrument throughput Characteristics of detectors (quantum efficiency,

noise)

SNRs in Astronomy

Fundamental limit set by photon statistics: SNR < √ N, where N = no. of detected source photons

Typical SNR's in Astronomy: Measures of astronomical EM fluxes:

Best precision: SNR ~ 1000 (0.1% error) Low by lab standards! Problems: difficulty of

calibration; faintness of interesting sources. Typical "good" measures: SNR ~ 20-30 Threshold detections: SNR ~ 5-10

Noise Sources Detector Noise (CCDs)

Read-noise (rms ~3-10e-1/read) Dark noise (3.10-4e-1/s/pixel)

Determined by Temperature of detector

Background Noise (Diffuse) Artificial light pollution Earth's atmosphere Ecliptic scattered sunlight Scattered Galactic light

Background Noise (Discrete) Exclusion zone around bright stars caused by scattered

light within instrument Source "confusion" caused by diffractive blending of

multiple faint sources

Magnitude System

An ancient and arcane, but compact and by now unchangeable, way of expressing brightnesses of astronomical sources. Magnitudes are a logarithmic measure of

spectral flux density (not flux!) Monochromatic Apparent Magnitudes

m = −2.5 log10 f − 21.1

where f is in units of erg.s−1.cm−2.Å−1

A system of “monochromatic magnitudes per unit wavelength”

Magnitude Normalization

Normalization is chosen to coincide with the zero point of the widely-used “visual” or standard “broad-band” V magnitude system: i.e. m(5500Å) = V

Zero Point: fluxes at 5500Å corresponding to m (5500Å) = 0, are (Bessell 1998)

f0 = 3.63.10−9 erg.s−1.cm−2.Å−1; or

fν0 = 3.63.10−20 erg.s−1.cm−2.Hz−1; or = 3630 Janskys

f0/hν = 1005 photons.s−1.cm−2.Å−1 is the

corresponding photon rate per unit wavelength

Surface Brightness

Surface Brightnesses (extended objects): μ = m + 2.5 log10

where m is the integrated magnitude of the source and is the angular area of the source in units of arcsec2.

1 arcsec2 = 2.35.10−11 steradians

μ is the magnitude corresponding to the mean flux in one arcsec2 of the source.

Absolute Magnitudes

M = m − 5 log10(D/10), where D is the distance to the source in Parsecs (pc) 1pc = 3.258 light-years or 3.086.1013

kilometers 1 pico-pc = a good day’s walk

M is the apparent magnitude the source would have if it were placed at a distance of 10 pc.

M is an intrinsic property of a source For the Sun, MV = 4.83

Source characterization Luminosity (L)

Power (energy/sec) radiated by source into 4 sterad Units: ergs.s−1

Flux (f) Power from source crossing normal to unit area at

specified location a distance D from source f = L/4D2 if source isotropic, no absorption Units: ergs.s−1.cm−2

Surface Brightness (I) Power per unit area per solid angle Units: ergs.s−1.cm−2.sterad−1 (f = I.)

I is independent of distance if source remains resolved

Point Source Sensitivity

Faintest UVOIR point sources detected: Naked eye:     5-6 mag Galileo telescope (1610):     8-9 mag Palomar 5m (1948):     21-22 mag (pg)

                                      25-26 mag (CCD) Keck 10m (1992):     27-28 mag HST (2.4m in space, 1990): 29-30 mag

NB: current optical detectors approach 100% QE ie: We can't improve sensitivity via detector

development. Improvements require new instrumentation.

Spatial Resolution Fundamental limit governed by diffraction in

telescope/instruments Min. image dia. (min)= 2.5/D rads(diffr. limit)

where D is the dia. of the telescope At 5500Å … min= 28”/D(cm) Inside Earth's atmosphere, turbulence strongly

affects image diameter. Resulting image blur & motion is called "seeing", and

typically yields: atm~0.7-1.5” i.e. spatial resolution in most instances is governed by

the atmosphere, not the telescope. Good site + Good environmental control Good AO – approaches diffraction limit

Spectral Resolution

Theoretical maximum governed by diffraction limts set by optical components: Practical limit set by photon rates High resolution devices are typically photon-

starved (except for Sun). ID's, surveys, classification at low resolution

10-500Å or 10 <R< 500 Physical analysis at moderate-to-high

resolution 0.1-10Å or 500 <R< 50,000

Highest to date: ~ 0.01Å or R = 500,000

Basic Lens formulae

Basic Mirror formulae

Refraction

2

1 n1

n2

Snell’s Law: n1 sin(1) = n2 sin(2)• n1 = refractive index in region 1• n2 = refractive index in region 2

where: n = c/v = vacuum /medium

Refraction andTotal Internal Reflection

ConstructiveInterference

DestructiveInterference

N=2 interfering beams

N >> 2 interfering Beams(eg: Grating Spectrograph)

Diffraction grating(N ~ 100,000)

Interferenceorder

wavelength

groove spacing incidence angle

diffracted angle

Michelson Interfermeter(N = 2 interference)

Fabry-Perot(N ~500)

Optics and Focus

Optics below represents a doublet lens Parallel rays from the left are made to

converge

Location where the rays cross is the “focal point”

Distance from the fiducial point in the lens is the “focal length” (fl)

fl

Images Object Plane Image Plane

These are “conjugates” of each other

Conjugate distances are: s1 ; s2

Lens formula 1/s1 + 1/s2 = 1/f

Magnification (m) is given by m = s1/s2

s1 s2

Object

Image

Focal length and focal-ratio (f/#)

Effective focal length (EFL = fl) is the distance from the optics to the focal point

f/# is the ratio focal length to the optic diameter (f/# = d/fl) f/1 is fast (v.difficult to control

aberrations) f/30 is slow (simple optics)fl

d

Plate Scale For a given optic with EFL = fl, the image

plane scale is given by: P.S. = 1c/fl (radians/m) = 206265 /fl (arcsec/mm)

For instance, a telescope with an EFL = 10m (eg: 1.2m @ f/8), plate scale is: 206265/104 = 20.6 arcsec/mm

However, a telescope with an EFL = 170m (eg: 10m @ f/17), plate scale is: 206265/1.7.105 = 1.2 arcsec/mm

ie: If you want a wide-field you have to have a small telescope

Entendu

Entendu = A (area * solid angle) Entendu is conserved for any optical

system ie: Conservation of Energy However, entendu can be lost in fibre systems

High entendu is a figure of merit for an optical system Equivalent to more energy (or information)

transport Telescope have generally to trade:

High A with High Spectrographs try to maximize A.R