Download D2 Stellar characteristics and stellar evolution

• The Electromagnetic Spectrum
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 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.
Wavelength and Color
• The color of light depends upon its wavelength.
𝐸 = ℎν
BLUE – HOT!!!
• SHORT Wavelength
• HIGH Frequency
• HIGH ENERGY!
RED – COLD!!!
• LONG Wavelength
• LOW Frequency
• LOW ENERGY!
Wien's law
Consequences of Wien's Law
Hot objects look blue.
Cold objects look red.
lpeak
T (ºK)
Sun
500 nm
5800
People
Neutron Star
9x103 nm
310
2.9x10-2 nm
108
Radio
10 m
0.03 K
Microwave
1 cm
3K
Infrared
1 mm
300 K
Visible
500 nm
6000 K
Ultraviolet
100 nm
100,000 K
X-Ray
0.1 nm
10 M K
The spectrum of
a
star/planet/object
reveals it’s
temperature
Hot!
Cool!
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.
Cooler
stars!
Hotter
stars!
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 requires 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!
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.
Key Features of the Atoms/Ions Spectra
Spectral lines
Hydrogen spectrum
The greater the difference between the quantum numbers, the
larger the energy of the photon emitted or absorbed.
Most
prominent
lines in many
astronomical
objects:
Balmer lines
of hydrogen
Analyzing Absorption an Emission Spectra
● Each element (atom/ion)
produces a specific set of
absorption (and emission) lines.
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)
● Comparing the relative
strengths of these sets of
lines, we can study the
composition of gases.
Step from line spectrum to continuous spectrum
The energy levels get closer together as the
quantum numbers get larger.
Key Features of the Continuum Spectra
► 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.
THERMAL
RADIATION
Hot, so it emits light
Peak color (wavelength) shifts to shorter
wavelengths as an object is heated
increasing temperature
So, emitted spectrum tells us about temperature
Temperature Spectrum of Objects
• All objects emit a continuous spectrum
• You are giving off light right now!
If you could fill a teaspoon just with material as dense as the
matter in an atomic nucleus, it would weigh ~ 2 billion tons!!
just for fun – so cute
Kirchoff's Laws of Spectroscopy/Radiation
Kirchoff formulated these laws empirically in the mid-19th
century – didn’t explain why – in early 20th century: QM –
nature of atom – beginning of understanding origin of spectra
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.
1. A hot solid, liquid or gas at high pressure produces a
continuous spectrum – all λ.
2. A hot, low-density / low pressure gas produces an
emission-line spectrum: energy only at specific λ.
3. A continuous spectrum source viewed through a cool,
low-density gas produces an absorption-line
spectrum: missing λ = dark lines.
Thus when we
see a spectrum
we can tell what
type of source we
are seeing.
Absorption Spectrum of Hydrogen Gas
Actual Examples of Emission Spectra
Element Spectrum
Argon
Helium
Mercury
Neon
Sodium
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.
What Color is Our 5800K Sun?
• The Sun emits all colors
• Blue-green is most intense
Peak Color (Wavelength) Depends
on Object’s Temperature
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.
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 CO2.
EX: (a) Explain the term black-body radiation.
Black-body radiation is that emitted by a theoretical perfect emitter for a given temperature.
It includes all wavelengths of electromagnetic waves from zero to infinity.
The diagram is a sketch graph of the black-body
radiation spectrum of a certain star.
(a) Copy the graph and label its horizontal axis.
(c) On your graph, sketch the black-body radiation
spectrum for a star that has a lower surface temperature
and lower apparent brightness than this star.
The red line intensity should be
consistently lower and the maximum
shown shifted to a longer wavelength.
(d) The star Betelgeuse in the Orion constellation emits radiation approximating to that
emitted by a black-body radiator with a maximum intensity at a wavelength of 0.97 µm.
Calculate the surface temperature of Betelgeuse.
The Amazing Power of Starlight
Just by analyzing the light received from a
star, astronomers can retrieve information
about a star’s
1. Total energy output
2. Surface temperature
3. Radius - radius
4. Chemical composition
5. Rotational and translational velocity relative to Earth
6. Rotation period
In general, the density and temperature of a star
decreases with distance from its centre. Because its
temperature is so high, a star’s core has to be
composed of high-pressure gases and not of molten
rock, unlike the cores of some planets
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.
Same total light emitted
Cool, but big.
There are a few stars that are
more luminous than the Sun.
For instance, Betelgeuse
has L ~ 14000 x Lsun.
There are lots
more low
luminosity stars
than high
luminosity stars.
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
Cepheid variables are stars with regular variation in luminosity (rapid brightening, gradual
dimming) which is caused by periodic expansion and contraction of outer surface (brighter as it
expands). This is to do with the balance between the nuclear and gravitational forces within the star.
In most stars these forces are balanced over long periods but in Cepheid variables they seem to take
turns, a bit like a mass bouncing up and down on a spring. Because they are so luminous it means that
very distant Cepheids can be observed from the Earth.
In 1784, the periodic pulsation of the supergiant star,
Delta Cephei, was discovered by the English amateur
astronomer, John Goodricke. This star has a period of
about 5.4 days
In 1908, the American astronomer, Henrietta Leavitt
working at the Harvard College Observatory compiled
a list of 1,777 periodic variables in the Large and
Small Magellanic Cloud . Eventually she classified 47
of these in the two clouds as Cepheid variables
published an article describing the linear relationship
between the luminosity and period of pulsation. Her
plot showed what is now known as the periodluminosity relationship:
Cepheids with longer periods are intrinsically more
luminous than those with shorter periods.
𝐿Θ 𝑖𝑠 𝑙𝑢𝑚𝑖𝑛𝑜𝑠𝑖𝑡𝑦 𝑜𝑓 𝑡ℎ𝑒 𝑠𝑢𝑛
It is now known that there are many more of these variable stars – collectively known as
Cepheids. The period of these stars varies between twelve hours and a hundred days.
Cepheid variable stars are known as “standard candles” because they allow us to
measure the distances to the galaxies containing Cepheid variable stars.
The distances, d, of the Cepheids can be calculated from the apparent brightness–
luminosity equation:
𝐿
𝑏=
4𝜋𝑑 2
The apparent brightness is measured using a telescope and CCD and the luminosity
is calculated from a measurement of the period of the Cepheid.
Cepheid stars are stars that have completed the hydrogen burning phase and moved off
the main sequence (see later for an explanation of this). The variation in luminosity
occurs because the outer layers within the star expand and contract periodically. This is
shown diagrammatically as:
(1) a layer loses hydrostatic equilibrium and is pulled
inwards by gravity
(2) the layer becomes compressed and less transparent to
radiation
(3) temperature inside the layer increases, building up the
internal pressure
(4) causing the layer to be pushed outwards
(5) During expansion the layer cools, becoming less dense
(6) and more transparent, allowing radiation to escape
and letting the pressure inside fall
(1) Subsequently the layer falls inwards under gravity
and the cycle repeats causing the pulsation of
the radiation emitted by the star.
EX: (a) Define (i) luminosity (ii) apparent brightness.
(i) Luminosity is the total power radiated by star.
(ii) Apparent brightness is the power from a star received by an observer on the
Earth per unit area of the observer’s instrument of observation.
(b) State the mechanism for the variation in the luminosity of the Cepheid variable.
Outer layers of the star expand and contract periodically due to interactions of
the elements in a layer with the radiation emitted.
The variation with time t, of the apparent brightness b, of a Cepheid variable is shown below.
Two points in the cycle of the star have
been marked A and B.
(c) (i) Assuming that the surface temperature of the star stays constant, deduce whether
the star has a larger radius after two days or after six days.
(i)
The radius is larger after two days (point A) because, at this time the
luminosity is higher and so the star’s surface area is larger.
(ii) Explain the importance of Cepheid variables for estimating distances to galaxies.
Cepheid variables show a regular relationship between period of variation of the luminosity
and the luminosity. By measuring the period the luminosity can be calculated and, by using
the equation b = L/4πd2 , the distances to the galaxy can be measured. This assumes that
the galaxy contains the Cepheid star.
(d) (i) The maximum luminosity of this Cepheid variable is 7.2 × 1029 W.
Use data from the graph to determine the distance of the Cepheid variable.
(ii) Cepheids are sometimes referred to as “standard candles”. Explain what is meant by this.
A standard candle is a light source of known luminosity. Measuring the period of a
Cepheid allows its luminosity to be estimated. From this, other stars in the same galaxy
can be compared to this known luminosity.
Let’s review some important things we want
to know about stars…
Given enough time and information, we can figure
out their…
• Brightness - easily observed
• Parallax to measure distance
• Spectral type - can get from the spectrum
• Brightness + Distance = Luminosity
• Temperature - can get from spectrum
• Temperature + distance = Size
• Mass - hard to figure out, but there are binary stars
• Age - exact age is hard, but can estimate
What do you do when you have
data and you don’t know what to do
with it and you don’t understand it?
CLASSIFY!
HOPE:
We just might get to know THE universe better?
Stars can be arranged into categories
based on the features in their spectra…
This is called
“Spectral Classification”
How do we categorize stars?
A few options:
1. by the “strength” (depth) of the absorption lines in their spectra
2. by their color as determined by their blackbody curve
3. by their temperature and luminosity
• Much of the work in classifying and explaining
stellar spectra and brightness was done by women
at Harvard around the turn of the century.
Harvard Computers (1890)
Annie Jump Cannon
(1863-1941)
• Single-handedly classified more
than 250,000 stellar spectra.
Henrietta Leavitt
(1868-1921)
Stars are classified
by their spectra as
O, B, A, F, G, K, and M
spectral types
• OBAFGKM
• hottest to coolest
• bluish to reddish
• An important sequence to remember:
–Oh Boy, An F Grade Kills Me
The Spectral Sequence
Class
Spectrum
Color
Temperature
O
ionized and neutral helium,
weakened hydrogen
bluish
31,000-49,000 K
B
A
F
neutral helium, stronger
hydrogen
blue-white
10,000-31,000 K
strong hydrogen, ionized
metals
white
7400-10,000 K
weaker hydrogen, ionized
metals
yellowish white
6000-7400 K
G
still weaker hydrogen, ionized
and neutral metals
yellowish
5300-6000 K
K
weak hydrogen, neutral
metals
orange
3900-5300 K
M
little or no hydrogen, neutral
metals, molecules
reddish
2200-3900 K
Eventually, the
connection was made
between the observables
and the theory.
Observable:
• Strength of Hydrogen Absorption Lines
• Blackbody Curve (Color)
Cecilia Payne
Theoretical:
• Using observables to determine
things we can’t measure:
Temperature and Luminosity
Categorizing the stars…
Hertzsprung-Russell (H-R ) Diagram
In the early 1900s, two astronomers, Ejnar
Hertzsprung in Denmark and Henry Norris Russell
in America, independently devised a pictorial way
of illustrating the different types of star.
• graph of luminosity versus temperature
(or spectral class)
Henry Norris Russell dissuaded Cecilia Payne-Gaposchkin from concluding that the
composition of the Sun is different from that of the Earth in her papers, as it
contradicted the accepted wisdom at the time. However, he changed his mind four
years later after deriving the same result by different means. After Payne was proven
correct, Russell briefly credited Payne for discovering that the Sun had a different
chemical composition from Earth in his paper. However the credit was still generally
given to him instead.[11]
Shematic H-R Diagram
O
B
A F G
K
M
BRIGHT
SUPER GIANTS
106
104
L/LΘ
GIANTS
102
10-2
FAINT
10-4
40,000
20,000
10,000
5,000
2,500
Temperature
HOT
COOL
the stars aren’t randomly scattered on this graph-- they form a line!Same
WHAT IS AMAZING: Stars of different masses fall along a narrow temperature,
path
in L/T diagram
but much
brighter than
MS stars
→ Must be
much larger
► Giant Stars
“Red Giants”
“Supergiants”
None of
these “extra”
stars are
Hydrogen
burning!
If a random star
falls on the Main
Sequence, you also
know that it’s
Hydrogen burning!
Nearly 90% of all
stars fit into this
category.
If you measure the
luminosity and the
color of a star, you
know its mass!!!
The more massive a star is, the more luminous it is…
But a higher rate of fusion means it’s burning its
fuel faster!
More massive stars are…
Low mass stars have
lifetimes comparable to the
Age of the Universe
High mass stars have very short
lifetimes, and disappear quickly!
•
•
•
•
Hotter
Brighter
Bigger
Shorter-lived
The high mass
stars are gone!
After time
passes…
Only long-lived
low mass stars
are left on the
main sequence!
Red Giants:
• Cool, but bright.
• Same temp as some
main sequence stars 
same surface
brightness!
Must be bigger AREA  BIGGER star!
(and thus the name, red giant)
Red giants are cooler than the Sun and so
emit less energy per square metre of
surface. However, they have a higher
luminosity, emitting up to 100 times more
L  T  4R
4
2
The Mystery of Red Giants and White Dwarfs…
Many of these stars have the same temperature as normal Main
Sequence stars, but they’re much brighter or fainter!
How is this possible???
Same Temperature & Surface Brightness
Hot.
Cool.
Same Luminosity
Cool, but big
Luminous!
If the size of the star changes,
its luminosity changes
Hot, but tiny
Faint.
4
2
L
=
b
x
Area
L  T  4R
The Hertzsprung–Russell (H – R) diagram is a scattergram of stars showing the relationship between the
stars' luminosities versus their surface temperature. It shows stars of different ages and in different
stages, all at the same time.
Vertical axis: luminosity/ luminosity of the Sun (L⊙= 3.839 × 1026 W). Axis is logarithmic and has no unit.
The temperature axis is also logarithmic and doubles with every division from right (low) to left (high).
main sequence stars: nearly 90% of all stars
▪ fusing hydrogen into helium, the difference between
them is in mass
▪ during the lifetime of a star its position will move on the
diagram as its temperature and luminosity changes
▪ left upper corner more massive than right lower corner
▪ cooler red stars relatively low luminosity;
▪ hotter blue stars: high luminosity.
Red giants are cooler than the Sun
• emit less energy/m2 of surface.
• higher luminosity means they have a much greater surface area.
• a much larger diameter than the Sun – “giant” stars.
White dwarfs
• remnants of old stars
• constitute about 9%of all stars
• energy not produced by nuclear fusion
• very hot when they stopped producing energy,
• they have a relatively low luminosity
• a small surface area.
• very small, hot, very dense stars
• take billions of years to cool down.
Supergiant stars are gigantic and very bright.
• A supergiant has 100 000 times the power and
at the same temperature of the Sun must have
a surface area 100 000 times larger: a diameter that
is over 300 times the diameter of the Sun.
• Only about 1%of stars are giants and supergiants.
• Relationship between the luminosity and the mass:
L ∝ M3.5 (Observations of thousands of main sequence stars)
L is the luminosity in W (or multiples of the Sun’s luminosity, L⊙ )
M is the mass in kg (or multiples of the Sun’s mass, M ⊙ ).
• Even a slight difference in the masses of stars results in a large difference in their luminosities.
• For a star to be stable it needs to be in hydrostatic equilibrium:
the pressure due to the gravitational attraction of inner shells = the thermal and radiation
pressure acting outwards.
For a stable star of higher mass there will be greater gravitational compression and so the core
temperature will be higher. Higher temperatures make the fusion between nuclei in the core
more probable giving a greater rate of nuclear reaction and emission of more energy;
thus increasing the luminosity.
• The mass of a star is fundamental to the star’s lifetime – High mass stars have shorter lifetimes.
– a star with a mass of 10 times the solar mass might only live 10 million years
compared to the expected lifetime of around 10 billion years for the Sun.
Formation of a star
▪ Gravitational attraction of hydrogen nuclei.
▪ Loss of PE leads to gain in KE and an increase in the gas temperature.
▪ The gas becomes denser and, when the protostar has sufficient mass, the temperature becomes high
enough for nuclear fusion to commence.
▪ The star moves onto the main sequence where it remains for as long as its hydrogen is being fused into
helium – this time occupies most of a star’s life.
▪ Eventually when most the hydrogen in the core has fused into helium the star moves off the main
sequence.
The fate of stars
▪ Star collapses when most of the hydrogen nuclei have fused into helium.
▪ Gravity now outweighs the radiation pressure and the star shrinks in size and heats up.
▪ The hydrogen in the layer surrounding the shrunken core is now able to fuse, raising the
temperature of the outer layers which makes them expand, forming a giant star.
▪ Fusion of the hydrogen adds more helium to the core which continues to shrink and heat up,
forming heavier elements including carbon and oxygen.
▪ The very massive stars will continue to undergo fusion until iron and nickel (the most stable
elements) are formed.
▪ What happens at this stage depends on the mass of the star.
A. Sun-like stars
▪ For stars up to about 4 solar masses the core temperature will not be high enough to allow
the fusion of carbon. This means that, when the helium is used up, the core will continue
to shrink while still emitting radiation.
▪ This “blows away” outer layers forming a planetary nebula around the star. When the
remnant of the core has shrunk to about the size of the Earth it consists of carbon and
oxygen ions surrounded by free electrons.
▪ It is prevented from further shrinking by an electron degeneracy pressure. Pauli’s exclusion principle
prevents two electrons from being in the same quantum state and this means that the electrons
provide a repulsive force that prevents gravity from further collapsing the star.
The star is left to cool over billions of years as a white dwarf.
Such stars are of very high density of about 109 kg m–3 .
The probable future for the
Sun is shown as the purple line
on the HR diagram.
B. Larger stars
▪ For the star much bigger than the Sun when in the red giant phase, the core is so large that
the resulting high temperature causes the fusion of nuclei to create elements heavier than carbon.
▪ The giant phase ends with the star having layers of elements with proton numbers that
decrease from the core to the outside (much like layers in an onion).
▪ The dense core causes gravitational contraction which, as for
lighter stars, is opposed by electron degeneracy pressure.
Even with this pressure, massive stars cannot stabilize.
▪The Chandrasekhar limit is the maximum mass of a stable
white dwarf. It is 1.4 times the mass of the Sun.
▪ When the mass of the core reaches this value the electrons
combine with protons to form neutrons – emitting neutrinos
in the process. The star collapses with neutrons coming as
close to each other as in a nucleus.
▪ The outer layers of the star rush in towards the core but
bounce off it in a huge explosion – a supernova.
▪ Although this process lasts just a few hours it results in the heavy
elements being formed.
▪ This blows off the outer layers and leaves the remnant core as
a neutron star.
▪ Now neutrons provide a neutron degeneracy pressure that resists
further gravitational collapse.
▪ The Oppenheimer–Volkoff limit places an upper value on a
neutron star for which neutron degeneracy is able to resist
further collapse into a black hole. This value is currently
estimated at between 1.5 and 3 solar masses.
PBS:
There is a thin line between a bang
and a whimper.
For stars, this line is called the
Chandrasekhar Limit, and it is the
difference between dying in a blaze
of glory and going out in a slow
fade to black. For our universe, this
line means much more: Only by
exceeding it can stars sow the seeds
of life throughout the cosmos.
Supernova 1987A with the
right hand image the region of the sky
taken just before the event
Black holes
Here we discuss their importance in astrophysics. It is not possible to form a neutron star having a
mass greater than the Oppenheimer–Volkoff limit instead the remnant of a supernova forms a
black hole. Nothing can escape from a black hole – including the fastest known particles, photons.
For this reason it is impossible to see a black hole directly but their existence can be strongly
inferred by the following.
● The X-rays emitted by matter spiralling towards the edge of a black hole and heating up. X-ray
space telescopes, such as NASA’s Chandra, have observed such characteristic radiation.
● Giant jets of matter have been observed to be emitted by the cores of some galaxies. It is
suggested that only spinning black holes are sufficiently powerful to produce such jets.
● The unimaginably strong gravitational fields have been seen to influence stars in the vicinity,
causing them to effectively spiral. A black hole has been detected in the centre of the Milky Way
and it has been suggested that there is a black hole at the centre of every galaxy.
EX: A partially completed Hertsprung–Russell (HR) diagram is shown below.
The line indicates the evolutionary path of the Sun from its
present position, S, to its final position, F. An intermediate
stage in the Sun’s evolution is labelled by I.
(a) State the condition for the Sun to move from position S.
Most of the Sun’s hydrogen has fused into helium.
(b) State and explain the change in the luminosity of the Sun that occurs between positions S and I.
Both the luminosity and the surface area increase as the Sun moves from S to I.
(c) Explain, by reference to the Chandrasekhar limit, why the final stage of the evolutionary
path of the Sun is at F.
White dwarfs are found in region F of the HR diagram. Main sequence stars that end up with a
mass under the Chandrasekhar limit of 1.4 solar masses will become white dwarfs.
(d) On the diagram, draw the evolutionary path of a main sequence star that has a mass of
30 solar masses.
The path must start on the main sequence above the Sun. This should lead to the super red giant
region above I and either stop there or curve downwards towards and below white dwarf in the
region between F and S.
First, there was a nebula.
The birth of stars in the M16 Eagle Nebula
Last, there was a nebula.
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 talking about hydrogen let me mention
one very interesting fact.
I call it: a WOW story of time!!!!!!
spin-flip transition even though the atom is in its ground state
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.
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).
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.
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.
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 cm3 of space.
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.