slide 1 telescope parameters light-gathering power (ability to see faint objects) resolving power...

Slide 1

Telescope parameters

• Light-gathering power (ability to see faint objects)

• Resolving power (ability to see fine details)

• Magnification (least important)

Slide 2

Other factors:

• Optical quality

• Atmospheric conditions

• Light pollution

Slide 3

SeeingWeather conditions and turbulence in the atmosphere set further limits to the quality of astronomical images.

Bad seeing Good seeing

Slide 4

The Best Location for a Telescope

Far away from civilization – to avoid light pollution

Slide 5

The Best Location for a Telescope (2)

On high mountain-tops – to avoid atmospheric turbulence and other weather effects

Paranal Observatory (ESO), Chile

Slide 6

Nick Suntzeff

Well-known astronomer at Cerro-Tololo, Chile

Now Physics Professor at TAMU

Cerro-Tololo observatory

Supernova in Centaurus A

Slide 7

Traditional Telescopes (1)

Traditional primary mirror: sturdy, heavy to avoid distortions.

Secondary mirror

Slide 8

Advances in Modern Telescope Design

2. Simpler, stronger mountings (“Alt-azimuth mountings”) to be controlled by computers

1. Lighter mirrors with lighter support structures, to be controlled dynamically by computers

Floppy mirror Segmented mirror

Modern computer technology has made possible significant advances in telescope design:

Slide 9

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

Slide 10

Slide 11

Examples of Modern Telescope Design (1)

Design of the Large Binocular

Telescope (LBT)

The Keck I telescope mirror

Slide 12

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.

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

Interferometry

Slide 13

Examples of Modern Telescope Design (2)

8.1-m mirror of the Gemini Telescopes

The Very Large Telescope (VLT)

Slide 14

Giant Magellan Telescope

Slide 15

CCD ImagingCCD = Charge-coupled device

• More sensitive than photographic plates• Data can be read directly into computer memory, allowing easy electronic manipulations

Negative image to enhance contrasts

False-color image to visualize brightness contours

Slide 16

The SpectrographUsing a prism (or a grating), light can be split up into different wavelengths (colors!) to produce a spectrum.

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

Slide 17

Exploring other wavelengths

• Radio

• Infrared

• UV

• X-ray

• Gamma-ray

Slide 18

Radio AstronomyRecall: Radio waves of ~ 1 cm – 1 m also penetrate the Earth’s atmosphere and can be

observed from the ground.

Slide 19

Science of Radio Astronomy

Radio astronomy reveals several features, not visible at other wavelengths:

• Neutral hydrogen clouds (which don’t emit any visible light), containing ~ 90 % of all the atoms in the Universe.

• Molecules (often located in dense clouds, where visible light is completely absorbed).

• Radio waves penetrate gas and dust clouds, so we can observe regions from which visible light is heavily absorbed.

Slide 20

Radio Telescopes

Large dish focuses the energy of radio waves onto a small receiver (antenna)

Amplified signals are stored in computers and converted into images, spectra, etc.

Slide 21

The Largest Radio Telescopes

The 100-m Green Bank Telescope in Green Bank, WVa.

The 300-m telescope in Arecibo, Puerto Rico

Slide 22

Radio InterferometryJust as for optical telescopes, the resolving power of a radio telescope is min = 1.22 /D.

For radio telescopes, this is a big problem: Radio waves are much longer than visible light

Use interferometry to improve resolution!

Slide 23

Radio Interferometry (2)The Very Large Array (VLA): 27 dishes are combined to simulate a large dish of 36 km in diameter.

Even larger arrays consist of dishes spread out over the entire U.S. (VLBA = Very Long Baseline Array) or even the whole Earth (VLBI = Very Long Baseline Interferometry)!

Slide 24

Very Long Baseline Interferometry

Slide 25

Radio observations with Very Long Baseline Interferometry (VLBI) are thousands of times more precise than optical observations (good enough to easily pinpoint a source the size of a pea in New York when sitting in Paris)

Slide 26

Slide 27

Frederick William Herschel 1738-1822

Discovery of the Infrared

He directed sunlight through a glass prism to create a spectrum (the rainbow created when light is divided into its colors) and then measured the temperature of each color. Herschel used three thermometers with blackened bulbs (to better absorb heat) and, for each color of the spectrum, placed one bulb in a visible color while the other two were placed beyond the spectrum as control samples. As he measured the individual temperatures of the violet, blue, green, yellow, orange, and red light, he noticed that all of the colors had temperatures higher than the controls. Moreover, he found that the temperatures of the colors increased from the violet to the red part of the spectrum. After noticing this pattern Herschel decided to measure the temperature just beyond the red portion of the spectrum in a region where no sunlight was visible. To his surprise, he found that this region had the highest temperature of all.

Slide 28

Cups with cold and hot water

chameleon

Slide 29

Infrared Astronomy ( ~ 1-300 m)

However, from high mountain tops or high-flying air planes, some infrared radiation can still be observed.

NASA infrared telescope on Mauna Kea, Hawaii

Most infrared radiation is absorbed in the lower atmosphere.

Slide 30

IRAS image of the Milky Way

Slide 31

NASA’s Space Infrared Telescope Facility (Now Spitzer Space Telescope)

Slide 32

Space Astronomy

Slide 33

The Hubble Space Telescope

• Avoids turbulence in the Earth’s atmosphere

• Extends imaging and spectroscopy to (invisible) infrared and ultraviolet

• Launched in 1990; maintained and upgraded by several space shuttle service missions throughout the 1990s and early 2000’s

Slide 34

Slide 35

Hubble Deep Field10 day exposure photo!

Over 1500 galaxies in a spot 1/30 the diameter of the Moon

Farthest and oldest objects are 12-13 billion ly away!

Space observations as a time machine

1011 galaxies in the observable universe

Slide 36

Ultraviolet Astronomy• Ultraviolet radiation with < 290 nm is

completely absorbed in the ozone layer of the atmosphere.

• Ultraviolet astronomy has to be done from satellites.

• Several successful ultraviolet astronomy satellites: IRAS, IUE, EUVE, FUSE

• Ultraviolet radiation traces hot (tens of thousands of degrees), moderately ionized gas in the Universe.

Slide 37

X-Ray Astronomy• X-rays are completely absorbed in the atmosphere.

• X-ray astronomy has to be done from satellites.

NASA’s Chandra X-ray Observatory

X-rays trace hot (million degrees), highly ionized gas in the Universe.

Slide 38

Gamma-Ray AstronomyGamma-rays: most energetic electromagnetic radiation;

traces the most violent processes in the Universe

The Compton Gamma-Ray Observatory

Slide 39

It takes 10,000 years for a photon emitted in the core to reach the surface!

Stars as black-body emitters

Slide 40

Black Body Radiation (1)The spectrum of a star’s light is approximately a thermal spectrum called a black body spectrum.

A perfect black body emitter would not reflect any radiation. Thus the name “black body”.

The spectrum of a black body emitter is described by a universal formula first suggested by Planck. It depends only on surface temperature.

Slide 41

Two Laws of Black Body Radiation

1. The peak of the black body spectrum shifts towards shorter wavelengths when the temperature increases. Wien’s displacement law:

max ≈ 3x106 nm / T(K)

(where T(K) is the temperature in Kelvin).

Slide 42

Color and Temperature

Orion

Betelgeuse

Rigel

Stars appear in different colors,

from blue (like Rigel)

via green / yellow (like our sun)

to red (like Betelgeuse).

These colors tell us about the star’s

temperature!

Slide 43

Two Laws of Black Body Radiation2. The hotter an object is, the more luminous it is.

= Stefan-Boltzmann constant

where A = surface area

L = A**T4

The Stefan-Boltzmann law:

sT 2

4

m

JFlux

Luminosity, or total power:

Radiation Flux, or power emitted by unit area of a black-body emitter, is proportional to the fourth power of its surface temperature:

428

K s m

J1067.5

Slide 44

K)(

nm103 6

T

Note units!!Wien’s law:

sT 2

4

m

JFlux The Stefan-Boltzmann law

Slide 45

Yellow light: ~ 520 nmMaximum of the black-body spectrum:

Example of black-body emitter: our sun

K)(

nm103 6

T

Surface temperature T =3x106 nm/520 nm 5800 K

Radius = 7x105 kmTotal radiated power (luminosity) L = T4 4R2 = 4x1026 W

The Stefan-Boltzmann laws

T 24

m

JFlux

Wien’s law

Slide 46

4

4

b

a

b

a

T

T

F

F

Comparing radiation fluxes and luminosities from two sources A and B:

42

b

a

b

a

b

a

T

T

R

R

L

L

Slide 47

The Spectra of StarsInner, dense layers of a

star produce a continuous (blackbody) spectrum.

Cooler surface layers absorb light at specific frequencies.

=> Spectra of stars are absorption spectra.

Slide 48

Fraunhofer in early 1800’s measures solar spectrum and identifies it with the spectrum of hydrogen in the lab

Bad news: stars are too far away to scoop their matter for testing

Good news: they consist of the same atoms as the stuff on the Earth

English astronomer Lockyer, in the late-1800's, discovered an unknown element in the Sun, i.e. a set of spectral lines which did not correspond to elements in the lab. He named this element helium (Latin for Sun element).

Slide 49

What is spectrum?

Slide 50

Light and MatterSpectra of stars are more complicated than pure blackbody spectra.

characteristic lines, called absorption lines.

To understand those lines, we need to understand atomic structure and the interactions between light and atoms.

Slide 51

Atom is mostly empty space!

Size of proton or neutron: ~10-15 m

Size of an electron cloud:~10-10 m (1 Angstrom)

Proton mass: 1.7x10-27 kgElectron mass: 9x10-31 kg

Slide 52

Thomson’s atom1899

Slide 53

Rutherford atom1911

Slide 54

Slide 55

“Planetary” model of atom

Proton mass: 1.7x10-27 kgElectron mass: 9x10-31 kg

Slide 56

Nuclear Density

If you could fill a teaspoon just with material as dense as the matter in an atomic nucleus, it would weigh ~ 2 billion tons!!

Neutron stars have such density

Slide 57

Different Kinds of Atoms• The kind of atom

depends on the number of protons in the nucleus.

Helium 4

Different numbers of neutrons ↔ different isotopes

• Most abundant: Hydrogen (H), with one proton (+ 1 electron).

• Next: Helium (He), with 2 protons (and 2 neutrons + 2 el.).

Slide 58

Slide 59

The atom contains a nucleus surrounded by a cloud of negatively charged electrons. The nucleus is composed of neutral neutrons and positively charged protons. The opposite charge of the electron and proton binds the atom together with electromagnetic forces.

Slide 60

•Matter is effected by forces or interactions (the terms are interchangeable) •there are four fundamental forces in the Universe:

•gravitation (between particles with mass) •electromagnetic (between particles with charge/magnetism) •strong nuclear force (between quarks) •weak nuclear force (that changes quark types)

Matter is effected by forces or interactions (the terms are interchangeable)

There are four fundamental forces in the Universe: gravitation (between particles with mass) electromagnetic (between particles with charge) strong nuclear force (between quarks) weak nuclear force (that changes quark types)

Slide 61

Catastrophe with atomsAccelerating electron produces EM radiation (light), loses energy and spirals into nucleus, i.e. atom should not work

Slide 62

Ultraviolet catastrophe with black-body radiation

Slide 63

There is a stable orbit (ground state) on which electrons do not radiate.

Changes of energy, such as the transition of an electron from one orbit to another around the nucleus of an atom, is done in discrete quanta. Quanta are not divisible. There is no ``in between''.

The quantization, or ``jumpiness'' of action as depicted in quantum physics differs sharply from classical physics which represented motion as smooth, continuous change.

Bohr’s atom

Slide 64

Atomic Transitions

• An electron can be kicked into a higher orbit when it absorbs a photon with exactly the right energy.

• All other photons pass by the atom unabsorbed.

Eph = E4 – E1

Eph = E3 – E1

(Remember that Eph = h*f)

Wrong energy

• The photon is absorbed, and

the electron is in an excited state.

Slide 65

Slide 66

secJ102

; 34

hh

kmvp

Perhaps one of the key questions when Bohr offered his quantized orbits as an explanation to the UV catastrophe and spectral lines is, why does an electron follow quantized orbits? The response to this question arrived from the Ph.D. thesis of Louis de Broglie in 1923. de Broglie argued that since light can display wave and particle properties, then perhaps matter can also be a particle and a wave too.

Energy and momentum of a particle are related to wavelength:

Wave-particle duality

Wave packetm

k

m

pE

22

222

mv

Slide 67

Your de Broglie wavelength:

msmkg

sJ

mv31

2

34

10/1010

10

de Broglie wavelength for the electron in an atom:

msmkg

sJ

mv10

529

34

10/1010

10

Note the velocity dependence!

Slide 68

The electron matter wave is both finite and unbounded. But only certain wavelengths will `fit' into an orbit. If the wavelength is longer or shorter, then the ends do not connect.

Thus, de Broglie explains the Bohr atom in that on certain orbits can exist to match the natural wavelength of the electron. If an electron is in some sense a wave, then in order to fit into an orbit around a nucleus, the size of the orbit must correspond to a whole number of wavelengths.

Why the orbits are quantized

Slide 69

If an electron is a wave around the atom, instead of a particle in orbit `where' is the electron at any particular moment?

The answer is that the electron can be anywhere around the atom. But 'where' is not evenly distributed. The electron as a wave has a maximum chance of being observed where the wave has the highest amplitude. Thus, the electron has the highest probability to exist at a certain orbit.

Werner Heisenberg

Erwin Shrödinger

1920s

Slide 70

Heisenberg’s Uncertainty Principle

px

tE

It is often stated that of all the theories proposed in this century, the silliest is quantum theory. Some say the only thing that quantum theory has going for it, in fact, is that it is unquestionably correct. - R. Feynman

Slide 71

Interference and diffraction of electron waves

Slide 72 p. 99

1859: Kirchhoff explains spectra of stars

Slide 73

Slide 74

Kirchhoff’s Laws of Radiation (1)1. A solid, liquid, or dense gas excited to emit

light will radiate at all wavelengths and thus produce a continuous spectrum.

Slide 75

Kirchhoff’s Laws of Radiation (2)2. A low-density gas excited to emit light will

do so at specific wavelengths and thus produce an emission spectrum.

Light excites electrons in atoms to higher energy states

Transition back to lower states emits light at specific frequencies

Slide 76

Kirchhoff’s Laws of Radiation (3)

3. If light comprising a continuous spectrum passes through a cool, low-density gas, the result will be an absorption spectrum.

Light excites electrons in atoms to higher energy states

Frequencies corresponding to the transition energies are absorbed from the continuous spectrum.

Slide 77

Analyzing Absorption Spectra• Each element produces a specific set of

absorption (and emission) lines.

By far the most abundant elements in the Universe

• Comparing the relative strengths of these sets of lines, we can study the composition of gases.

Slide 78

Lines of HydrogenMost prominent lines in many astronomical objects: Balmer lines of hydrogen

Slide 79

The Balmer Linesn = 1

n = 2

n = 4

n = 5n = 3

H H H

The only hydrogen lines in the visible wavelength range.

Transitions from 2nd to higher levels of hydrogen

2nd to 3rd level = H (Balmer alpha line)2nd to 4th level = H (Balmer beta line)

…

Slide 80

Observations of the H-Alpha LineEmission nebula, dominated by the red H line.

Slide 81

Absorption Spectrum Dominated by Balmer Lines

Modern spectra are usually recorded digitally and

represented as plots of intensity vs. wavelength

Slide 82

The Balmer ThermometerBalmer line strength is sensitive to temperature:

Almost all hydrogen atoms in the ground state (electrons in

the n = 1 orbit) => few transitions from n = 2 => weak

Balmer lines

Most hydrogen atoms are ionized => weak Balmer

lines

Slide 83

Measuring the Temperatures of Stars

Comparing line strengths, we can measure a star’s surface temperature!

Slide 84

Spectral Classification of Stars (1)

Tem

pera

ture

Different types of stars show different characteristic sets of absorption lines.

Slide 85

Spectral Classification of Stars (2)

Mnemonics to remember the spectral sequence:

Oh Oh Only

Be Boy, Bad

A An Astronomers

Fine F Forget

Girl/Guy Grade Generally

Kiss Kills Known

Me Me Mnemonics

Slide 86

Stellar Spectra

OB

A

F

GKM

Surface tem

perature