Transcript Slide 1

Star Evolution
Version: 5/4/08
Assembled by
Ken Mitchell
Livermore TOPScience
dtd. 3/15/09
Star Forming Region NGC 3582
See Notes on next slide
What's happening in the NGC 3582 nebula?
Bright stars and interesting molecules are forming. The complex
nebula resides in the star forming region called RCW 57.
Visible in this image are dense knots of dark interstellar dust, bright
stars that have formed in the past few million years, fields of
glowing hydrogen gas ionized by these stars, and great loops of
gas expelled by dying stars.
A recent detailed study of NGC 3582 uncovered at least 33 massive
stars in the end stages of formation, and the clear presence of the
complex carbon molecules known as polycyclic aromatic
hydrocarbons (PAHs).
PAHs are thought to be created in the cooling gas of star forming
regions, and their development in the Sun's formation nebula five
billion years ago may have been an important step in the
development of life on Earth.
Stellar evolution looks at the life history of stars
By what we understand now, stars shine with the energy they produce
through nuclear fusion.
The particulars of the fusion process in a given star may change during
the course of its life, and produce spectacular changes in the appearance
and makeup of the star.
The types of changes that stars undergo, depending on their size, is the
topic of stellar evolution.
Stars make the universe interesting. Basically every chemical element
that we know of, except hydrogen and helium, comes from the variety of
nuclear reactions that occur in the late phases of a massive star's life.
Further, it is primarily by the explosions of massive, dying stars that
these newly created elements are injected into the interstellar medium.
From there, they can be incorporated into new stars and planets.
Hertzsprung-Russell Diagram
The Luminosity Formula
L=4πR2σT4
Where:
L = Luminosity
R = Star’s radius
σ = Stefan-Boltzmann constant
= 5.669 x 10-5 erg/cm2 x deg4 x sec
T = temperature (OKelvin)
Summary: The Importance of Mass
Here's what the mass of a star can tell us:
1. Its main sequence properties: radius,
luminosity, temperature.
1. How long it can spend on the main sequence
(MS).
3. What kind of fuels it can “burn” after leaving
the MS.
4. What kind of stellar remnant it will end up as.
Carina Nebula where a maelstrom of a star's birth
and death is taking place. See Notes
The Birth of Stars
Stars form from the nebulae and debris of massive stars
that supernovae at the end of their life. This provides
building materials for the next generation of stars,
Gravity and shock-compression of gases cause
Hydrogen and heavier elements to accrete around a mass
to form new stars. These are called protostars.
The star-forming process may also create planets in its
accretion disk.
The types of stars formed (and their lifetime) will depend
on the available material that can be accreted.
Accretion Disk
A rotating disk of gas
and dust matter that may
form around any of a
variety of stars or other
massive objects.
In the case of young
stars, accretion disks
contain unconsolidated
material, such as cosmic
dust grains, which may
subsequently accrete to
form planets and other
sizable objects
See Notes
The Beginning – A Protostar
A star's life is a constant battle against gravity, the force that wants
to compress it, and pressure support which is trying to hold it up.
That battle begins as soon as a molecular cloud begins to collapse
into a protostar.
Main Sequence Star
During the main sequence lifetime, the star achieves stability by
gas pressure (P=nkT); the collapse of the star is prevented
because if the star were to compress slightly, it would get hotter
and the rate of nuclear fusion would increase, causing it to expand
again.
Leaving the Main Sequence: Red Giants
Star in the Process of Dying See Notes
Death of Low Mass Stars: White Dwarfs – 1.
All this time the core continues to collapse and heat up, until
finally it becomes hot enough to fuse helium into carbon by a
reaction called the triple alpha process. Carbon can also fuse
with helium to form oxygen.
At this stage, the contraction of the core stops: the star has a
helium-burning core surrounded by a hydrogen-burning shell.
Eventually, of course, even the helium in the core becomes
exhausted -- this happens after only about 100 million years
because helium burning is less efficient than hydrogen burning.
Now the core is mostly carbon and oxygen (C-O) and once again
continues its collapse.
The star produces energy by burning helium in a shell around the
core and hydrogen in a shell around that shell.
White Dwarf – 2.
How dense does the C-O core
become? About 106 g (1 ton)
per cubic centimeter! This hot,
superdense core is called a
white dwarf. In the final
stages of a low-mass star's
life, the outer layers of the
star become unstable and get
ejected into space, forming a
planetary nebula.
Fusion in the hydrogen and helium shells stops and the white
dwarf becomes exposed, slowly cooling because it is no longer
generating energy. Since degeneracy pressure is unaffected by
temperature, the white dwarf does not contract as it cools; it
basically cools off like a charcoal briquette, eventually becoming a
cold black dwarf. See Notes
Death of Low-Mass Stars: White Dwarfs – 3.
Stars smaller than about 5 solar masses (low mass stars) never get hot
enough to get fusion in their C-O cores at the end of their red giant phase.
The collapse of the core therefore continues until a new kind of pressure
is able to support the star against gravity. This pressure is called electron
degeneracy pressure.
Electron degeneracy pressure is a consequence of an important principle in
quantum mechanics called the Pauli exclusion principle. It basically says
the following: no two electrons in a material can have exactly the same
quantum coordinates.
If you try to squeeze a lot of electrons in a small volume, they are forced
into a smaller range of space-type quantum coordinates. This makes them
be in higher and higher quantum energy levels, so they will tend to move at
higher and higher “velocities.”
This has nothing to do with temperature: even at absolute zero (0 K), the
electrons must continue whizzing around, because the Pauli exclusion
principle forbids them from coming to rest.
This constant motion of electrons produces a pressure which increases
with density; thus, it only becomes important in very dense matter.
The Chandrasekhar Limit
The Chandrasekhar (chan-druh-SAY-kar) limit is the maximum
nonrotating mass which can be supported against
gravitational collapse by electron degeneracy pressure. It is
commonly given as being about 1.4 solar masses.
Computed values for the limit will vary depending on the
approximations used, the nuclear composition of the mass,
and the temperature.
As white dwarf stars are supported by electron degeneracy
pressure, this is an upper limit for the mass of a white dwarf.
Main-sequence stars with a mass exceeding approximately 8
solar masses therefore cannot lose enough mass to form a
stable white dwarf at the end of their lives, and instead form
either a neutron star or black hole.
The Triple Alpha Process
While beryllium-8 is present, its creation is a small energy sink. To
release energy, carbon-12 and heavier elements must be created.
Carbon-12 is created when helium-4 combines with beryllium-8. In
this interaction, carbon-12 nucleus is left in an energetic state from
which it decays, releasing a gamma-ray. The conversion of beryllium-8
into carbon-12 releases 7.37 MeV.
The conversion of helium-4 into carbon-12 is therefore
accomplished through the following two reactions:
He4 + He4→Be8 (2.6 × 10-16 sec lifetime)
Be8 + He4→C12 + γ
The process of converting three helium-4 nuclei into a single carbon12 nucleus releases a total of 7.27 MeV, all of which remains trapped
within the star. This fusion chain can be treated as a single process;
it is then called the triple-alpha process (an alpha particle is a helium4 nucleus).
Death of High Mass Stars: Neutron Stars -- 1.
Stars more massive than about 5 solar masses don't stop
with helium burning.
Hydrostatic equilibrium predicts that for a star supported by
gas pressure (P=nkT), the larger the mass, the higher the
temperature at the center of the star must be to support that
mass.
So massive stars have higher temperatures in their cores.
We also know that the reason we need high temperatures to
get fusion is that positively charged nuclei will repel away
from each other before they can fuse unless they're moving
very fast. It stands to reason, therefore, that if you're trying
to fuse nuclei with more protons, you need higher
temperatures to do it.
That is, the more massive a star is, the heavier the nuclei
which it can fuse.
Death of High Mass Stars: Neutron Stars -- 2.
Thus, as helium is used up in the core and a C-O core develops,
that core does NOT collapse all the way to become a white dwarf.
Before doing so, it becomes hot enough to fuse carbon into neon,
and oxygen into sulfur and silicon. Finally, silicon gets fused into
iron.
Every time a heavier element is made, it sinks to the center of the
star where it eventually becomes hot enough to undergo fusion.
The ash of yesterday's burning becomes the fuel for today's.
The result is an ``onion-skin'' structure for the star, with:
1. hydrogen on the outside, and inner shells of
2. helium,
3. carbon/oxygen,
4. oxygen/neon/magnesium,
5. sulfur/silicon, and
6. an iron core at the very center.
Death of High Mass Stars: Neutron Stars – 3.
Fusion occurs in all of the shells simultaneously, but can't occur in the core
because you can't get energy from fusing iron into heavier elements. So the
iron core cannot generate energy and begins to shrink.
In less than a day, silicon burning in the dying star produces so much iron
that the iron core exceeds the Chandrasekhar limit. Because the core is not
generating its own energy, it is supported only by electron degeneracy
pressure, and once the mass of the core exceeds 1.4 solar masses it must
collapse. In the process of the collapse, many of the electrons in the core
are squeezed so tightly with the nuclei that they merge with protons to
become neutrons.
This reduces the electron degeneracy pressure further and accelerates the
collapse. In less than a second, the core collapses into a ball of neutrons
only about 10 km in radius -- a neutron star. A neutron star is supported by
gravity against a different kind of pressure, neutron degeneracy pressure.
This is analogous to electron degeneracy pressure but requires much
greater densities to become important. In fact, a neutron star is so dense
that a cubic centimeter of it weighs as much as all the people on the planet
Earth put together – (which is??)
Types of Supernovae
Supernovae are divided into two basic physical types:
Type Ia. These result from some binary star systems in which a
carbon-oxygen white dwarf is accreting matter from a
companion. (What kind of companion star is best suited to
produce Type Ia supernovae is hotly debated.) In a popular
scenario, so much mass piles up on the white dwarf that its
core reaches a critical density of 2 x 109 g/cm3. This is enough
to result in an uncontrolled fusion of carbon and oxygen, thus
detonating the star.
Type II. These supernovae occur at the end of a massive star's
lifetime, when its nuclear fuel is exhausted and it is no longer
supported by the release of nuclear energy.
If the star's iron core is massive enough then it will collapse
and become a supernova.
Type II Supernovae
Type II Supernovae – 1.
This sudden collapse of a massive star's core into a volume
over a million times smaller than its original volume is really
bad news for the star.
The outer layers of the star come raining down onto the
core. Somehow this collapse changes into an explosion: a
Type II supernova.
The process by which this happens is still being
investigated, but evidently the core collapses to something
below its equilibrium radius and then rebounds slightly.
That bounce transfers an enormous amount of energy to
the layers falling down from above.
Ultraviolet flash of light produced from a dying star
just before it exploded. See Notes
Type II Supernovae – 2.
A strong wave of energy -- a shock wave -- travels out through
the envelope and heats the star so much that the outer layers
are blown away.
Another important effect is the huge numbers of neutrinos that
are produced when the neutron star is formed. Ordinarily,
neutrinos don't interact much with matter, but these neutrinos
are so numerous and energetic that they help push the outer
layers of the star away.
The total amount of energy released in a Type II supernova is
about 1053 ergs. About 99% of that energy is emitted as
neutrinos, whereas only 1% is converted into the kinetic and
heat energy of the ejecta (i.e., outer gas layers). Yet enough
light is emitted by a supernova to make it as bright as a billion
Suns.
Type II Supernovae – 3.
In the process of a supernova explosion, the temperatures are
briefly so high (billions of OK) that elements heavier than iron can
be produced.
Remember, we can't get energy by fusing an iron nucleus with
another nucleus, but who's to say we can't provide enough energy
to make this happen?
That energy needed to make elements beyond iron is readily
available during the supernova explosion. Unfortunately, many of
these elements are hard to detect with spectroscopy so
observational proof is still lacking.
We do know, however, that many or most of the elements beyond
iron had to have been created very rapidly, so supernovae are still
the best bet.
The remnant of
a supernova ...
stars that
exploded 9
billion years
ago have led to
new insights
into dark
energy.
Type II Supernovae – 4.
But the significance of supernovae goes far beyond the
production of rare elements. Even light nuclei like carbon
and oxygen, which can be produced by low mass stars,
would be locked up inside white dwarfs if it wasn't for
supernovae.
Supernovae enrich the gas between the stars with all kinds
of chemical elements that are necessary for the production
of planets and life.
The pressure of a supernova blast may trigger the formation
of stars and planets in an interstellar cloud of gas and dust.
Supernovae are also thought to be the major source of high
energy cosmic rays, which can affect the evolution of life by
causing mutations.
Type II Supernovae – 5.
The most important supernova that has happened in modern
astronomical history is known as SN 1987A, and became visible,
as the name suggests, on February 24, 1987. The explosion
occurred in the Large Magellanic Cloud.
By studying the spectrum and the apparent brightness of SN
1987A, astronomers confirmed many of the ideas for how Type II
supernovae occur.
They even had pictures of the star before it exploded. It was a
blue supergiant star with a mass of around 20 solar masses and a
luminosity of around 5 solar masses. They found evidence of
radioactive 56Co in the SN's spectrum. (This isotope of cobalt is
radioactive with a short half-life, indicating that it was freshly
synthesized in the star.)
Experiments on Earth, which look for neutrinos from the Sun,
witnessed a sudden burst of neutrinos just before the SN became
visible, supporting another theoretical prediction.
Pictures of Supernova 1987A -- “Before" and “After"
See Notes
Summary of Type II Supernovae – 6.
To summarize, we learned from SN 1987A that:
Type II supernovae represent the deaths of massive stars, as
expected.
The supernova indeed releases lots of heavy elements into the
interstellar medium.
Neutrinos are produced in copious amounts, as the theories
predicted.
Pulsars: Stellar Beacons
What's left after a Type II supernova?
Pulsars: Stellar Beacons – 1.
If theories are correct, the collapsed core should remain as a
neutron star. Like a white dwarf, a neutron star is thought to
have a maximum mass before it becomes unstable.
That mass is around 3 solar masses.
A neutron star below this limit should just sit around and
cool off like a white dwarf.
But neutron stars are so tiny that even a very hot one would
have a very low luminosity (L=4πR2σT4 ).
So how could you possibly observe one?
Pulsars: Stellar Beacons – 2.
In 1967, a graduate student in England named Jocelyn Bell was
looking at data from a radio telescope and found, much to her
surprise, that one radio source was emitting a pulse of radiation
every 1.33 seconds.
For a while, astronomers thought this might be a signal from
intelligent extraterrestrials: “little green men.” Soon, however,
many more of these sources were found, including one in the Crab
supernova remnant that emits 30 pulses a second.
Astronomers now believe that these sources, dubbed pulsars, are
rapidly spinning neutron stars with strong magnetic fields. The
rotating magnetic field produces an electric current, in much the
same way that an electric generator operates here on Earth.
As the electrons in the current are accelerated, they emit
electromagnetic radiation in a sort of conical beam. The radiation
can have a broad range of wavelengths, from radio waves through
X-rays. Each time that beam sweeps by us, we see a burst of
radiation (akin to the flashing of a lighthouse beacon).
Illustration of a
Pulsar
Pulsars: Stellar Beacons – 3.
How do we know that the pulsars are neutron stars? The fact that
there's a pulsar in the Crab supernova remnant is suggestive,
because we'd expect a neutron star to be left behind. Also, the
Crab pulsar is slowing down at a rate that is consistent with an
age of about 1000 years (plus light travel time), which is when the
supernova was observed.
The sharpness of the pulses indicates that they come from a
region roughly 100 km across, since otherwise the finite travel
time for light to move across the region would smear out the
pulse.
The masses of pulsars in binary systems can be calculated (from
the observed or deduced motions of the stars) and give masses
between 1.4 and 1.8 solar masses, too massive for a white dwarf
but just right for a neutron star.
Finally, all other reasonable explanations, such as oscillation or
rotation of white dwarfs, do not explain the rapidity and regularity
of the pulses.
Pulsars: Stellar Beacons – 4.
Why do pulsars have strong magnetic fields? When the iron
core collapses to become a neutron star, the magnetic field lines
remain trapped inside it -- they get scrunched up. This increases
the strength of the magnetic field by a huge factor (1010 to 1012).
What provides the energy for a pulsar? The root source is the
rapid rotation of the neutron star, which results from the collapse
of a much larger, rotating core. Again, the idea is that when an
object collapses, the angular momentum must stay constant, so
the object must spin faster.
As the pulsar radiates energy, then, we expect its rotation to
gradually slow down. Eventually the rotation is too slow to
produce much radiation and the “lighthouse beacon” gradually
fades from view.
By carefully monitoring pulsars over periods of months, we can
learn more about the interiors of neutron stars.
Why the Pulsar is Flashing
Pulsar Geometry from Rotating Neutron Star
Black Holes: Gravity's Ultimate Victory – 1.
I mentioned that the maximum possible mass of a neutron star
was about 3 solar masses. This is analogous to the
Chandrasekhar limit, but remember that the Chandra limit
applies to white dwarfs and is about 1.4 solar mass.
What happens if a star is so massive that the supernova
explosion is unable to blow enough matter away to keep the
neutron star below this limit?
Such a star might have an original mass of, say, 50 solar
masses or more. It turns out that there is no known source of
pressure that can support the core against gravitational
collapse once it exceeds about 3 solar masses.
When the core becomes so small that the escape velocity at its
surface exceeds the speed of light, it becomes a black hole.
Jets from a Black Hole
An artist's conception of
the blazar BL Lacertae as
it spurts out jets of
charged particles
accelerated by corkscrew
magnetic field lines.
Note: Copy URL into browser
for site to get movie and
explanation. (~90MB)
http://tinyurl.com/556af7
Examining the Throat of a Black-Hole Jet – 2.
Black-hole jets are among the most violent, and important, phenomena in
the universe. When matter falls toward a black hole, it often forms an
extremely hot "accretion disk" around the hole as it spirals in. Somehow,
despite the intense gravity hauling things inward, the inner part of the
accretion disk often manages to emit narrow jets of matter from its top and
bottom faces. Typically, a few percent of the infalling stuff ends up getting
expelled this way, powered somehow by the energy of the rest falling in.
These accretion-disk jets can extend millions of light-years from the cores
of active radio galaxies and quasars. We also see them in miniature
squirting from "microquasars" in our own galaxy, when enough matter
feeds into a stellar-mass black hole.
The spinning accretion disk develops an intensely powerful magnetic field.
(Because the gas is very hot, electrons get stripped off atoms and can
move around freely, so the gas is electrically conducting — a plasma — a
lot like the copper conductors in a spinning generator.) The disk's spin
wraps up the magnetic field into a tight spiral; the most intense, tightly
wound part of the field breaks loose and squirts upward from the disk's
inner poles; and plasma (because it's tied to magnetic fields, just like any
moving conductor is) flies out along with it.
Black Holes: Gravity's Ultimate Victory – 3.
A black hole has no “surface” but does have a size scale
associated with it, the Schwarzschild radius. Sometimes called
the event horizon, this describes a sphere around the black hole
separating objects which can (in theory) escape the black hole's
gravity from those which cannot. At the event horizon, the escape
velocity is the speed of light (c); inside the event horizon, not even
light can escape the black hole. The Schwarzschild radius is
given by: Rsch = 2GM/c2 where M is the mass of the black
hole and G is Newton's gravitational constant = 6.668 x 10-8 dyne
2
2
cm /gm . For the Sun, this radius is only 3 km.
Bear in mind that well outside the event horizon, the black hole
behaves just like any other object of the same mass. So if the Sun
suddenly became a black hole, the Earth would continue orbiting
it just as before.
Schwarzschild Radius
The Schwarzschild (SHWORTS-sheeld) radius is a
radius determined for a massive object by setting
the escape velocity equal to the speed of light.
For black holes, this radius is larger than the object
itself, creating an event horizon from which
nothing, not even light, can escape.
(It was named for Karl Schwarzschild, who first
calculated it.)
Black Holes: Gravity's Ultimate Victory – 4.
As the formula above shows, the event horizon of a black hole
grows larger as it accumulates more mass.
In the centers of many galaxies, we believe there are black
holes that have accumulated so much mass (from eating
interstellar gas and stars) that they have masses of up to 109
solar masses.
These super-massive black holes may have quite different
origins from the stellar black holes that result from individual
massive stars. However, even for one of these “monsters” the
corresponding Schwarzschild radius is only about 20 A.U. -much smaller than our solar system.
So the universe is in no danger of being eaten up by black
holes.
Illustration of a "blazar" showing jets of plasma
See Notes
Black Holes: Gravity's Ultimate Victory – 5.
How do we observe black holes? Although the black hole itself
emits no light, it can make itself known via its gravity. If we see a
star orbiting an unseen companion that is well over 3 solar
masses in mass, and if there is evidence that the companion is
very small, we can be pretty sure it's a black hole.
One indicator of small size is the emission of X-rays from hot gas
falling in towards the companion. (X-rays have very high
energies, but these energies would be impossible just outside a
black hole's event horizon because the orbital speeds there
approach the speed of light.)
Internal friction of the gas falling towards the black hole could
then account for the X-ray emission.
In our Galaxy, there are a handful of bright X-ray sources which
are likely candidates for stellar black holes.