The study of the universe is so challenging, astronomers cannot
ignore any source of information; that is why they use the entire
spectrum, from gamma rays to radio waves. In astronomy, we cannot
perform experiments with our objects (stars, galaxies, ). The only
way to investigate them, is by analyzing the light (and other
radiation) which we observe from them. The Electromagnetic
Spectrum
Slide 2
Wavelength and Color The color of light depends upon its
wavelength. BLUE HOT!!! SHORT Wavelength HIGH Frequency HIGH
ENERGY! RED COLD!!! LONG Wavelength LOW Frequency LOW ENERGY! E =
hf
Slide 3
The Amazing Power of Starlight Just by analyzing the light
received from a star, astronomers can retrieve information about a
stars 1.Total energy output 2.Surface temperature 3.Radius
4.Chemical composition 5.Velocity relative to Earth 6.Rotation
period
Slide 4
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
stars temperature.
Slide 5
Hot! Cool! Hotter stars! Cooler stars! Through a telescope the
'star' Albireo is revealed to actually be composed of two bright
stars. These stars have very different surface temperatures, as
indicated by their color.
Slide 6
We only think of visible wavelengths as being light, but a lot
of familiar things involve light of different wavelengths! Size of
atomic nucleus! Size of Mt. Everest!
Slide 7
Where Does Light Come From? The following had been known during
the 19th century: accelerated charges emit energy and hence produce
light If we picture an electron as in orbit around the nucleus, it
should radiate light changing direction is acceleration! (a force
is required from something to change direction) This caused a major
problem with classical physics If the electron radiated due to its
motion around the nucleus, it would lose energy and soon spiral
into the nucleus. The world should collapse instantly!
Slide 8
1913, Niels Bohr formulated 3 rules regarding atoms:
1.Electrons can only be in discrete orbits. 2. A photon can be
emitted or absorbed by an atom only when an electron jumps from one
orbit to another. 3. The photon energy equals the energy difference
between the orbits. The discrete (quantum) nature of the energy
"levels" of the electron gives Quantum Mechanics its name.
Slide 9
The greater the difference between the quantum numbers, the
larger the energy of the photon emitted or absorbed. Key Features
of the Atoms/Ions Spectra Spectral lines Hydrogen spectrum Most
prominent lines in many astronomical objects: Balmer lines of
hydrogen
Slide 10
Observations of the H-Alpha Line Emission nebula, dominated by
the red H line.
Slide 11
Hydrogen and helium account for nearly all the nuclear matter
in today's universe. The abundance of hydrogen by mass, is 73%
(Helium is 25%, All other elements 2%). By atomic abundance,
hydrogen is 90%, helium 9%, and all other elements 1%. That is,
hydrogen is the major constituent of the universe. As we are
already talkig about hydrogen let me mention one very interesting
fact. I call it: a WOW story of time!!!!!! When a proton captures
an electron to form hydrogen atom (after Big Bang) the spins of the
two particles either point in the same direction or opposite
direction. Due to the spin both electrons and protons behave like
little magnets. Hydrogen atom with antiparallel spins is more
stable (the electron is more tightly bound to the proton) than the
atom with the two spins parallel because unlike poles attract each
other and like poles repel. Since a anti-parallel-spin capture is
three times as probable as an parallel-spin capture, 75% hydrogen
atoms in between the stars is with spins anti-parallel and 25% with
parallel-spin. spin-flip transition even though the atom is in its
ground state
Slide 12
This means that the electron in the hydrogen atom with its spin
parallel to that of the proton will ultimately flip over so that
its spin is anti-parallel to the proton spin emitting 21-cm photon
(radio part of the spectrum). 11,000,000 years every few hundred
years We should not ordinarily expect to receive much 21-cm
radiation because it takes on the average about 11,000,000 years
for an electron in an unmolested hydrogen atom to flip over to an
anti-parallel spin position. But there are so many such atoms in
interstellar space that the occasional emission of a 21-cm photon
from any one of them adds up to the observed intensity of the 21-cm
line. The collisions among the neutral hydrogen atoms that occur
every few hundred years keep the number of hydrogen atoms with
parallel and antiparallel electrons spins constant or in thermal
equilibrium. This hydrogen radio line with a wavelength of 21 cm
was first predicted theoretically in 1944 by the Dutch astronomer
H. C. van de Hulst (highly unlikely to be seen in a laboratory on
Earth). It was observed in radio telescopes at Harvard and then
throughout the world in the early 1950s.
Slide 13
Most of what is known about the distribution of cold gas in the
Galaxy, including the mapping of the nearby spiral arms, has come
from detailed studies of the variation of 21-cm line of Hydrogen
emission across the sky. Stars radiate all freq. but visible light
won't penetrate the dust clouds and this 21 cm will giving us
informations. Reading the intensity of 21 cm hydrogen line from
different parts of Universe we can get a fairly reliable picture of
the distribution of gas (hydrogen) and dust throughout the galaxy.
From the data that have now been collected, we know that the dust
in interstellar space, which constitutes only about 1% of the
interstellar material, is almost entirely responsible for the
dimming of the stars. The size of a dust grain is about 0.000001 cm
(about the size of the wavelength of visible light) and only one
such grain, on the average, is present in each 10,000,000,000 cm 3
of space.
Slide 14
Analyzing Absorption an Emission Spectra Comparing the relative
strengths of these sets of lines, we can study the composition of
gases. atom/ion Each element (atom/ion) produces a specific set of
absorption (and emission) lines. spectral signature fingerprints We
call this the "spectral signature" or fingerprints of an atom/ion.
Allows the identification of elements across the galaxy and
universe. (If we mapped it and can recognize it)
Slide 15
The energy levels get closer together as the quantum numbers
get larger. Step from line spectrum to continuous spectrum
Slide 16
If an electron is given enough energy (via a photon or by other
means) it can escape the atom. The electron is then "unbound" and
the quantization of energy levels disappears. The energy of an
electron in the continuum is not quantized. A hot, dense object
contains many "loose" electrons which can emit photons of any
energy. The light produced by a hot, dense object is called thermal
emission and it contains photons of all energies, i.e. light of all
colors, or wavelengths. The resulting "rainbow" is called a
continuous spectrum. As we heat up an object, we are giving the
electrons more kinetic energy, so they become able to emit more
energy. The hotter the object becomes,the brighter the continuous
spectrum becomes. Key Features of the Continuum Spectra
Slide 17
THERMAL RADIATION Hot, so it emits light
Slide 18
increasing temperature So, emitted spectrum tells us about
temperature Peak color (wavelength) shifts to shorter wavelengths
as an object is heated
Slide 19
Temperature Spectrum of Objects All objects emit a continuous
spectrum You are giving off light right now!
Slide 20
If you could fill a teaspoon just with material as dense as the
matter in an atomic nucleus, it would weigh ~ 2 billion tons!!
Kirchoff's Laws of Spectroscopy/Radiation Kirchoff formulated these
laws empirically in the mid-19th century didnt explain why in early
20 th century: QM nature of atom beginning of understanding origin
of spectra just for fun so cute Kirchhoff did not know about the
existence of energy levels in atoms. The existence of discrete
spectral lines was later explained by the Bohr model of the atom,
which helped lead to quantum mechanics.Bohr model quantum
mechanics
Slide 21
1. A hot solid, liquid or gas at high pressure produces a
continuous spectrum all . 2. A hot, low-density / low pressure gas
produces an emission-line spectrum energy only at specific . 3. A
continuous spectrum source viewed through a cool, low-density gas
produces an absorption-line spectrum missing dark lines. Thus when
we see a spectrum we can tell what type of source we are
seeing.
Slide 22
Absorption Spectrum of Hydrogen Gas
Slide 23
Actual Examples of Emission Spectra ElementSpectrum Argon
Helium Mercury Neon Sodium
Slide 24
The Spectra of Stars Inner, 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 25
To understand the nature, to interpret many beautiful phenomena
you have to have a tool. We are introducing something that we know
all about and then well compare the nature with that ideal
case!!!!!!! A black body is a theoretical object that absorbs 100%
of the radiation that hits it. Therefore it reflects no radiation
and appears perfectly black. A black body is a theoretical object
that absorbs 100% of the radiation that hits it. Therefore it
reflects no radiation and appears perfectly black. In practice no
material has been found to absorb all incoming radiation, but
carbon in its graphite form absorbs all but about 3%. It is also a
perfect emitter of radiation. At a particular temperature the black
body would emit the maximum amount of energy possible for that
temperature. This value is known as the black body radiation. It
would emit at every wavelength of light as it must be able to
absorb every wavelength to be sure of absorbing all incoming
radiation. Blackbody Radiation
Slide 26
In 1900 Max Planck characterized the light coming from a
blackbody. The equation that predicts the radiation of a blackbody
at different temperatures is known as Planck's Law. Note that the
peak shifts with temperature. The peak emission from the blackbody
moves to shorter wavelengths as the temperature increases (Wiens
law) The hotter the blackbody the more energy emitted per unite
area at all wavelengths. Bigger object emit more radiation.
Slide 27
Wien's Law The wavelength of the maximum emission of a
blackbody is given by: peak T (K) Sun500 nm5800 People9x10 3 nm310
Neutron Star 2.9x10 -2 nm10 8 Radio 10 m0.03 K Microwave 1 cm3 K
Infrared 1 mm300 K Visible 500 nm6000 K Ultraviolet 100 nm100,000 K
X-Ray 0.1 nm10 M K The spectrum of a star reveals its
temperature
Slide 28
Consequences of Wien's Law Hot objects look blue. Cold objects
look red. Except for their surfaces, stars behave as
blackbodies.
Slide 29
An interesting example of a black body radiation is the thermal
emission of the Earth (or any other body). This thermal emission
(also called infrared emission, due to its characteristic
wavelength) is due to the Earth's temperature. In the picture, we
can see the real emission compared to a black body radiation of a
body at a temperature of 280K. In the picture it is also possible
to see the absorption spectral lines of oxygen and CO 2.
Slide 30
What Color is Our 5800K Sun? The Sun emits all colors
Blue-green is most intense Peak Color (Wavelength) Depends on
Objects Temperature
Slide 31
Thermal radiation can explain much of this spectrum Thermal
radiation at 225K Thermal radiation at 6000K This is a spectrum of
Mars! The 6000K radiation is reflected light from the sun. The 225K
radiation is thermal emission from the planet.
Slide 32
But what is that other stuff? Emission lines! Extra light at
very specific wavelengths. Absorption lines! Light has been removed
at very specific wavelengths.
Slide 33
Sometimes emission lines dominate the light output. The Cygnus
supernova remnant emits almost all of its light as emission lines!
PS. A supernova remnant is the expanding shell of hot gas left over
after a star explodes.
Slide 34
A caveat to the rules of thermal radiation! Cooler objects can
sometimes emit more light overall. Decreasing temp means less light
emitted per unit area. An object can compensate by being BIGGER.
Lower surface brightness, but larger surface area. Hot. Cool. Cool,
but big. Same total light emitted
Slide 35
The Luminosity of a Star How bright a star or galaxy really is
is described by its luminosity. The energy radiated by a star is
emitted uniformly in all direction. If we regard stars to be black
body radiators, the luminosity L of a star is given by
Stefan-Boltzmann law for a black body: L = 4 R 2 T 4 (W) (all )
Doubling the temperature increases the luminosity by a factor of
16. A hotter star is more luminous than a cooler one of the same
radius. A bigger star is more luminous than a smaller one of the
same temperature. The total energy emitted by the star per second
(i.e. power) is called Luminosity of the star, L. The luminosity of
a 100 W lightbulb is (approximately) 100 W if you measure over all
wavelengths. (Most of it is in the infrared; the part of the
luminosity in visible wavelenghts is less.)
Slide 36
d By the time the energy arrives at the Earth it will be spread
out over a sphere of radius d (distance). apparent brightness
energy received per unit time per unit area (W/m 2 ) radiation
sensitive instrument If d can be measured then the luminosity of
the star can be determened. This is very important property to know
as it gives clues to the internal structure of the star, its age
and its future evolution. Since we can't go to a star to measure
its luminosity, we have to be clever. If we know the distance to
the star we can do it, because there is a simple relation between
the distance d to the star, the apparent brightness b of the star,
and the luminosity L of the star. It is easy to measure the
apparent brightness of a star, a galaxy, a supernova,... by
bolometer attached to a telescope.
Slide 37
A person radiates ~ 100 W = 100 J/s So that the energy output
in a day is I dont know exactly how to use this information to lose
extra weight that I gained recently.
Slide 38
How can we measure the Sun's luminosity? We can measure the
intensity of sunlight at the Earth. This should include all
wavelenghts. b ~ 1.4 x 10 3 W/m 2. So we should multiply b by the
area of this sphere: L = 4 d 2 b = 4 x 3.14 x (1.5 x 10 11 m) 2 =
2.8 x 10 23 m 2 So the luminosity is (1.4 x 10 3 W/m 2 ) x (2.8 x
10 23 m 2 ) 3.90 x 10 26 W. distance Sun-Earth can be measured
using different method Sometimes youll find luminosity of our Sun
written as: L = 3.90 x 10 26 W
Slide 39
Betelgeuse ("beetle juice"), a red supergiant star about 600
light years distant - one of the brightest stars in the familiar
constellation of Orion, the Hunter - the first direct picture of
the surface of a star other than the Sun. While Betelgeuse is
cooler than the Sun, it is more massive and over 1000 times larger.
If placed at the center of our Solar System, it would extend past
the orbit of Jupiter (has an immense but highly variable, outer
atmosphere ). As a massive red supergiant, it is nearing the end of
its life and will soon become a supernova.brightest starsOrion, the
HunterpictureBetelgeusecooler than the Sunour Solar Systemsupernova
There are a few stars that are more luminous than the Sun. For
instance, Betelgeuse has L ~ 14000 x L sun. There are lots more low
luminosity stars than high luminosity stars.
Slide 40
However: knowing how bright a star looks doesnt really tell us
anything about the star itself! We cannot see their size! if we
anly knew the distance
Slide 41
From Brightness to Luminosity If we know the distance to a
star, we can calculate the luminosity (energy output) Parallax is
the most direct measure of distance It is very hard to measure the
distance.
Slide 42
We know how big the Earths orbit is, we measure the shift
(parallax), and then we get the distance July distance of the star
from the earth, d January Parallax angle, p D=1 AU background stars
of the total angular shift. it does not make any significant
difference which distance you want to talk about from sun or from
earth For small angles: or click me
Slide 43
D = 1 AU = 149597870691 m 1.5x10 11 m when talking about stars,
parallax is very, very small number. Parallaxes are expressed in
seconds. so 1 pc = 3.09 X 10 16 m Parsec is short for parallax
arcsecond Parsec: distance where a star shifts by 1 over a 1/2
year
Slide 44
Parallax has its limits The farther away an object gets, the
smaller its shift. Eventually, the shift is too small to see. Even
the nearest star has a tiny parallax! First measured in 1838 The
closest bright star Alpha Centauri 4.3 light-year 0.75 pc In 1989,
the European Space Agency (ESA) launched a satellite called
Hipparcos to accurately measure the positions and motions of nearly
120,000 stars - plus about another million or so stars with good,
but lower precision.
Slide 45
Parallaxes give us distances to stars up to perhaps a few
thousand light years. Beyond that distance, parallaxes are so small
than they cannot be measured with contemporary instruments.
Astronomers use more indirect methods beyond a few thousand light
years.
Slide 46
Since both are red (the same color), the spectra peak at the
same wavelength. By Wien's law By our law governing Luminosity,
radius, and temperature of an object (star!) Star A is 9 times
brighter and as they are the same distance away from Earth star A
is 9 times more luminous: So, Star A is three times bigger than
star B. Suppose I observe with my telescope two red stars A and B
that are part of a binary star system. Star A is 9 times brighter
than star B. What can we say about their relative sizes and
temperatures? then they both have the same temperature. L = 4 R 2 T
4 (W) It must be that star A is bigger in size (since it is the
same temperature but 9 times more luminous). How much?
Slide 47
Suppose I observe with my telescope two stars, C and D, that
form a binary star pair. Star C has a spectral peak at 350 nm -
deep violet Star D has a spectral peak at 700 nm - deep red What
are the temperatures of the stars? By Wien's law Thus we have for
star C, and for star D
Slide 48
If both stars are equally bright (which means in this case they
have equal luminosities since the stars are part of a pair the same
distance away), what are the relative sizes of stars C and D? So
that stars C is 4 times smaller than star D.
Slide 49
The Doppler Effect observed change in wavelength of an EM wave
due to relative velocity betwee the source and the detector The
light of a moving source is blue/red shifted by / 0 = v r /c
Wavelength change v r = radial velocity 0 = actual wavelength
emitted by the source
Slide 50
of EM waves coming from stars are very often greater than those
obtained in the laboratory emitted from same elements (He, H).
Redshift shift toward greater wavelength Our Sun supercluster of
galaxies, BAS11 v=0.07c d = 1 billion LY Based on calculations
using the Doppler effect, it appears that nearby galaxies are
moving away from us at speed of about 250,000 m/s. The distant
galaxies are moving away at speeds up to 90 percent the speed of
light. The universe is moving apart and expanding in all
directions. Edwin Hubble in the 1920's. The more distant a galaxy,
the faster it is moving away. Some galaxies, however, are moving
toward us, and their light shows a blue shift.
