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Astronomy 101
Lecture 18, Mar. 31, 2003
Observing night
The observing session using our 14 inch telescope is tomorrow (Tuesday)
evening (Wednesday if cloudy on Tuesday). The forecast is not good, so
check the course web page
http://www.astro.sunysb.edu/dpeterso/AST101/index.html
for cancellation notice. Session will start at 7:00 PM to be able to view
Saturn. Come to the roof of the 4th floor of the ESS building.
Second Midterm
The exam is next Monday, April 7, in lecture. Same format as Midterm 1.
Exam will cover lectures 11 through 18 (chapters 16 through 21). We will have
some time for review in Wednesday’s lecture. Send requests for topics for
review to me -- [email protected]
Astronomy 101
Lecture 18, Mar. 31, 2003
Evolution of massive stars and stellar explosions – (Chapter 21 in text)
Old white dwarf stars never die, they just fade away ??
Well, not exactly. If the white dwarf is part of a binary pair and can attract
matter (mostly hydrogen) from its companion, the matter spirals into the very hot
surface of the white dwarf.
The hot hydrogen deposit undergoes nuclear burning, leading to a very bright
burst of light, X rays, etc. that lasts for about 2 months with maximum
luminosity 104 times normal – Nova (new star)
Nova
The rapid consumption of the infalling hydrogen by the usual pp
chain and CNO cycle burning causes a brief flare of light. The
increase in luminosity can be a factor of 10,000 (decrease of 10
magnitudes.
This can happen repeatedly as more matter falls on the white dwarf,
so some novae have been seen at the same star many times.
L
The white dwarf also slowly grows in mass as matter from its companion
falls on it – often with spectacular later consequences !
Recall that the pressure inside a white dwarf (a hot ball of carbon) is
mainly due to degenerate electrons – the unfriendly electrons that don’t
like to get close to each other.
When the mass of the white dwarf member of the binary grows to 1.4
solar masses (called the Chandrasekhar limit), the weight of the outer
layers is now greater than what can be sustained by degenerate electron
pressure. Without the ability for the pressure to counteract gravity, the
weight of the star pushes the star inward rapidly.
A runaway collapse then occurs, compressing the core and raising
its temperature to the point that most of the carbon can ignite.
Carbon burning occurs throughout most of the star, leading to a
cataclysmic explosion that rips the star apart and roughly a
billion-fold increase in light over the period of a month or two.
Type Ia supernova
Type I Supernova
Ordinary evolution of star
through supergiant, planetary
nebula explusion and white
dwarf formation
Accretion of matter from
companion on white dwarf until
reach 1.4 solar masses; then
supernova explosion.
Type I supernovae (actually type Ia) that occur by accretion of material on a
white dwarf surface release a well defined amount of energy since they all
explode from the same initial mass star, so the peak luminosities of all Type
Ia supernovae are about the same. Knowing that all Type Ia Supernovae rise
to the same luminosity makes it a STANDARD CANDLE.
Shape of light
curve –
luminosity vs.
time – tells us
what type the
supernova is.
Type I
Other kinds of Type I supernovae formed by collisions of two white
dwarfs do not have a fixed luminosity.
Use of the Standard Candle.
use L = 4pR2 I0
Where I0 is the intensity at the supernova
surface.
At a distance d from the supernova, the
intensity Iapp is lower by
Iapp/I0 = R2/d2
So,
I0 R2 = Iapp d2
And thus
(inverse square law)
(cross multiply)
L = 4pd2 Iapp
Measure Iapp (at earth), know L because it is Ia supernova
(by solving this equation for d) .
calculate d
Supernovae are so bright (10 billion times luminosity of sun) that they can
be seen in very distant galaxies. The distance to the supernova tells how
far away the galaxy is.
Evolution of massive stars – differs from solar
mass stars
Stars heavier than about 8 solar masses
develop sufficiently high core temperatures
to burn the carbon ash from helium burning.
There are many reactions possible now for further
burning – for example 12C + 12C → 24Mg + energy and
4He + 12C → 16O + energy. Similar 4He capture
reactions can lead all the way up to formation of iron:
56Fe. The core temperatures reach a few billion
degrees K at this stage.
One of the paths from silicon to nickel (
56Ni
56Ni
) by 4He reactions
is unstable and decays to
56Fe
In these late burning phases, there are shells of burning all the way from
hydrogen on the outside to iron formation at the center. The burning of
silicon to iron takes place in less than a day because the core temperatures
have now risen to a few billion degrees (MS sun is 15 million degrees).
But Iron is the end of the line !
There is no way to arrange protons and neutrons into a collection of nuclei
that is lighter than Fe – that is to say, no reactions exist to liberate energy
from iron.
A massive star that has evolved to an iron center quits core burning; gravity
now wins over the pressure and starts squeezing the core – raising the
temperature toward 10 billion degrees.
Mass of iron nucleus is less
than the sum of smaller
nuclei, so small nuclei ‘burn’
by fusion into iron and
release energy.
Nuclei above iron have less
mass than sum of smaller
nuclei it can break into, so
such nuclei as lead and
uranium can ‘burn’ by fission –
they split into lighter nuclei.
Type II Supernovae
When the massive star iron core contracts to have T = 1010 K, it is so hot
that its blackbody radiation is very intense and has very short
wavelengths (lT = constant : Wien’s Law). Short wavelength means high
energy:
lf = c ;
E = hf
(c and h are constants; f is frequency)
The photons (gamma rays) have energy large enough that they can tear
apart nuclei:
e.g.
g + 56Fe → 32S + 24Cr
g + 32S →
16O
+
16O
and so forth, until the nuclei are all broken down into protons & neutrons,
in a bath of electrons and photons.
These photodisintegration reactions take place in less than 1 second once
they start. They eat up energy (absorb) energy. (Just the reverse of
the burning reactions that powered the MS star).
The loss of energy causes Temperature to go down, reducing the
pressure and gravity further collapses the star.
Type II Supernovae
When the collapsed core reaches a density of 1012 kg/m3, the
electrons and protons are very tightly packed and very energetic. The
reaction
p + e- → n + n
becomes possible.
The neutrinos escape the star altogether, causing even more energy
loss and a consequent loss of pressure. Collapse is now essentially free
fall.
When the density increases to 1015 kg/m3, the neutrons in the core are
so close together that they become degenerate (neutrons, like
electrons are ‘unfriendly’ and they resist additional squeezing).
When the star develops this neutron degeneracy condition, it is like
hitting a brick wall. The outer falling layers hit the wall of the very
very high neutron pressure and rebound outward. The shock wave
radiating outward blows the outer layers of the star off in a massive
Supernova explosion. (Type II supernova – implosion of massive star)
The expanding disk from the ejection of the very hot material causes a
tremendous growth in luminosity – up to 1010 times the luminosity of the
sun, visible over vast reaches of space.
before
after
after
In 1987, a Type II supernova was observed in the Large Magellenic Cloud
– the galaxy nearest our Milky Way. SN1987a was observed on earth in
the visible light and in neutrinos, giving much information about the
supernova collapse. The star before the collapse was found on older
pictures so its properties were known.
Neutrinos made by the p + e → n + n reaction just before the explosion,
were observed in huge underground detectors in the US and Japan over
a several second period. More energy escaped with the neutrinos than
in light!
The Crab ‘nebula’
In 1054, Chinese astronomers recorded a
supernova in our galaxy that lit up the
daytime sky. Today, we see the rapidly
expanding filaments of gas from the site
of the explosion (now a neutron star). The
expanding material radiates in the visible
light and X-ray portion of the spectrum.
Measuring the velocities (Doppler effect)
we can see that the material was all
concentrated at the star in 1054.
Type I (carbon detonation
in accreting white dwarfs)
and Type II (iron core
collapse) supernovae reach
similar peak luminosities.
Type II are not standard
candles since the light
output depends on the mass
of the collapsing star which
varies.
Since the Type I supernovae are explosions of mainly carbon filled white
dwarfs, there is very little hydrogen and helium seen in the spectrum. Type
II supernovae blew off large amounts of hydrogen and helium and the
spectra show these elements. (presence/absence of hydrogen and helium
lines tells us immediately whether Type I or II.
Supernovae are very rare events – only about one in our galaxy each century.
But they are so bright that we can observe them in distant galaxies.
Formation of the elements
The universe started out containing mainly hydrogen (90%) and helium
(about 10%), formed at the birth of the universe. The nuclear burning in
stars transforms some of this material into heavier elements – carbon,
oxygen, calcium etc. up to iron. When a supernova erupts, these newly
formed elements are spewed out into space and are incorporated into
subsequent stars, planets – and life !
The elements observed from stellar
spectra are indicated in the graph. (Note
the vertical scale is ‘logarithmic’ – each
tick mark is a factor of 100 in number of
atoms present). The observed abundances
agree quite well with the expectations
based on studies of nuclear reactions in
our labs, and on the known stability of the
various nuclei – up to iron.
The elements heavier than iron are NOT
formed from burning in stellar cores (iron
is the most stable nucleus) !
Formation of elements beyond iron
OK – so nuclear burning and supernova distribution puts elements up to iron
into play. But how do we get the elements like copper and gold that are
heavier than iron?
First answer: in the final stages of burning in the interiors of evolved stars,
both iron and neutrons (byproducts of burning reactions) are present. Can get
reactions that build up the heavier elements. For example:
n+
56Fe
→
then
n + 57Fe →
and
n+
But
59Fe
decays to
58Fe
59Co
59Fe
57Fe
58Fe
→ 59Fe
(stable nucleus)
(stable)
(unstable)
Fe has 26 protons,
Co has 27
→ 59Co + e- + n producing a stable element higher
(more protons) in the periodic table.
This build up of heavy elements takes time, since it waits on a free neutron to
wander along and initiate the next reaction. But after time, it allows the
formation of elements as heavy as 209Bi (Bismuth) with 83 protons. Beyond
that, the slow process (s-process) fails since a new nucleus formed by neutron
capture decays back to Bi before the time a new neutron happens by.
Formation of elements beyond bismuth
The elements beyond Bismuth (like thorium, uranium, plutonium) occur by the
rapid process (r-process) during supernovae explosions. In the first minutes
of a supernova, many free neutrons from the nuclear reactions are present.
With the very large density of neutrons, the chain of elements beyond Bi can
be formed before the unstable nuclei decay.
Direct evidence for formation of
radioactive nuclei come from the
supernova light curves.
The radioactive nuclei 56Ni and 56Co are
formed from neutron reactions on 56Fe in
the explosion. The electrons from their
decay generate light as they are bent in
the intense magnetic fields. We see the
light from a supernova declines with time
with two components – one with the half
life measured on earth for 56Ni and the
longer one appropriate for 56Co.
Thus stars are the furnaces in which all the elements necessary for life
(carbon, oxygen, nitrogen) are made. Moreover all the elements upon which
our technological civilization rests (iron, silicon, gold, lead, uranium etc.)
are also forged in the furnaces of stellar interiors or in supernovae
explosions.
The material ejected by the supernovae seed the interstellar gas clouds
from which subsequent stars will be born. Thus the very old populations of
stars (the first generation stars) are deficient in the heavy elements,
while the youngest stars show the largest concentrations of heavy
elements.
Star formation
Insterstellar
medium
Stellar
recycling
Supernovae
Stellar evolution