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