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Number 124
ISSN 0110-7801
NZ
iron in stars
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Iron is only one of almost 100 elements of the periodic
table that are found in the Sun and other stars, but it is
one of the most important elements in how it affects the
evolution of a star. Iron is the last element that is fused at
the core of the star before it disintegrates in a supernova.
Since no more energy can be gained to hold up the star,
the star undergoes gravitational collapse. The reason
no more energy can be gained comes down to how
the fundamental components of the nucleus combine
together.
How is the nucleus held together?
If you measure the mass of an atomic nucleus and
then add up the mass of the all its constituent protons
and neutrons, you will find that they do not have the
same value. The nucleus has less mass than the parts
that make it up. If you convert this mass difference to
an energy using E=mc2, this value is called the binding
energy. It can be thought of as a measure of how much
energy is required to break apart the nucleus. Dividing
this binding energy by the number of nucleons (the
number of protons plus the number of neutrons) gives
the binding energy per nucleon. This number is the key
reason why iron is so important to the life of stars (see
Figure 1).
Average binding energy per nucleon (MeV)
9
8
7
6
5
4
3
2
1
0
0
30
60
90
120
150
180
210
240
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Number of nucleons in nucleus
Figure 1: As you change the number of nucleons in a
nucleus, there is also a change in the binding energy per
nucleon. Due to the interplay between the strong nuclear
force and electromagnetic force there is a peak in this
energy around iron-56.
As you move through the periodic table from hydrogen
to sodium (the stable isotope of which has a nucleus
containing 23 nucleons – 11 protons and 12 neutrons),
there is an increase in the binding energy per nucleon.
As more nucleons are added to the nucleus, the
combined strong nuclear force packs the nucleus tighter,
making it harder to break apart. The strong nuclear force
is a fundamental force that keeps the positively charged
protons in the nucleus together.
From magnesium (24 nucleons) to xenon (131 nucleons)
there is a relatively flat section. This feature is associated
with the growing importance of the electromagnetic
repulsion between the protons. The strong nuclear force
and the electromagnetic force are able to keep each
other in balance, such that there is no real change in
the energy needed to break apart the nucleus. Beyond
caesium (134 nucleons), there is a slow decline as the
nucleons on one side of the nucleus cannot exert a
strong enough nuclear force on the nucleons on the
other side but the electromagnetic repulsion is being
felt.
It is this curve that tells us why a fission reactor uses
uranium and a fusion power plant would use hydrogen.
To the left of the peak, atoms become more tightly
bound if they are combined to form a single nucleus
that lies nearer to the peak. This process is called fusion.
For atoms to the right of the peak, they would be more
tightly bound by breaking into atoms that are nearer to
the peak. This process is fission.
It is during the flat section from magnesium to xenon
that there is a peak at around iron. The peak actually falls
at nickel-62, not iron-56 as is commonly stated. In fact,
iron-58 also has a higher binding energy per nucleon
(Fewell). The reason why iron-56 is more important than
these elements to the life cycle of some stars comes
down to helium nucleus.
iron
Iron is the last element that is fused at the core
of the star before it disintegrates in a supernova,
as Jeffrey Simpson, post graduate student in the
Department of Physics and Astronomy at the
University of Canterbury, explains:
How do stars evolve and explode?
Stars shine as a result of the release of energy due to
nuclear fusion. Their evolution is primarily controlled by
their mass.
In a massive star (5 times the mass of the Sun or more),
the evolution will follow a different track to that of
the Sun. This is due to its size which allows for higher
energies and pressures to be obtained at its core where
the fusion takes place.
After it has exhausted its fuel of hydrogen at the core, by
fusing it to helium, the core contracts under gravitational
pressure. This pulls hydrogen down to deeper layers so
that a shell of hydrogen fusion forms around the nonburning helium core. This shell of hydrogen then causes
the outer layers to expand due to the increased thermal
energy, producing a red giant. Eventually the pressure
and temperature at the core reaches a level that helium
fusion can take through what is known as the triplealpha process. Three helium nuclei are fused to form one
carbon-12.
At a temperature of 200 million kelvin, a carbon atom
can fuse with another helium nuclei or alpha particle to
form oxygen-16. As the star ages, higher temperatures
are reached as one source of fuel is used up and the
star’s core contracts. Over time the conditions are
reached such that more alpha particles can be captured
to form neon-20, silicon-28, and then finally nickel-56
and iron-56 (see Figure 2).
The star is now like an onion (see Figure 3), with layers
of different materials tracing its chemical evolution. At
its core is iron and nickel. As stated above, nickel-62 is
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Figure 2: Carbon burning. Two carbon atoms fuse to form a neon-20 nucleus plus a helium nucleus and some energy in the
form of a gamma ray.
Diagram adapted by Jeffrey Simpson from: http://en.wikipedia.org/wiki/File:CNO_Cycle.svg
actually the most stable element in terms of its nuclear
binding energy per nucleon. But it is difficult to create
in the interior of stars through the addition of abundant
species. It is almost certainly produced in stars, but not in
great enough quantities to compete with iron in terms of
being relevant to stellar evolution.
The iron/nickel core mass increases until it reaches
the Chandrasekhar limit, which is about 1.4 times the
mass of the Sun. At this mass, electron degeneracy
pressure is overcome. This is the pressure caused by
the Pauli Exclusion Principle which forbids fermions
(such as electrons) being in the same energy state. It is
an extremely strong force but can be overcome. At this
point, electrons and protons combine to form
neutrons and the core collapses, leading to a
supernova.
There are several different types of supernovae. In terms
of our story of iron in astronomy, it is the type Ib, Ic and
Figure 3: This diagram shows a simplified (not to scale)
cross-section of a massive, evolved star (with a mass
greater than eight times the Sun). Where the pressure and
temperature permit, concentric shells of Hydrogen (H),
Helium (He), Carbon (C), Neon/Magnesium (Ne), Oxygen
(O) and Silicon (Si) plasma are burning inside the star.
Ref: http://commons.wikimedia.org/wiki/File:Evolved_star_fusion_shells.svg
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II supernovae that we are interested in. These are the
supernovae where iron plays a key role.
So iron having nearly the highest binding energy per
nucleon is what causes a core collapse supernova such
as the one that caused SN 1054, creating a source of light
so bright it could be seen during the day. This was the
progenitor of the Crab Nebula (see Figure 4).
Another example of a core-collapse supernova was SN
1987A, which occurred in the Large Magellanic Cloud
(Figure 5). One of its discoverers was a New Zealander,
amateur astronomer Albert Jones who has made
500,000 observations of variable stars.
The other type of supernova is Type 1a (Figure 6).
This is caused by a white dwarf star exceeding the
Chandrasekhar limit. White dwarfs are the remnants of
stars like our Sun which do not explode in supernova.
Instead, they do not have any more fusion at their core
beyond carbon and oxygen. Their outer layers are lost
and the core is left naked in space. In type 1a supernova,
the white dwarf gains matter from a companion star
which allows it to exceed the Chandrasekhar limit. Here
iron is produced, but this is explosively in the core of
the star. It is also produced by the radioactive decay of
nickel-56 to cobalt-56 and then iron-56.
Figure 4: The Crab Nebula as seen by the Hubble Space
Telescope in 2000. Japanese and Chinese astronomers
recorded this violent event nearly 1,000 years ago in 1054.
It is believed to be the result of a type II supernova.
Ref: http://hubblesite.org/newscenter/archive/releases/2005/37/)
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iron
Figure 5: This Hubble Space Telescope image, taken in
February 1994, with the Wide Field and Planetary Camera
2, shows the full system of three rings of glowing gas
surrounding supernova 1987A.
Courtesy of P. Challis (Harvard-Smithsonian Center for Astrophysics
Are there stars with no iron or other heavy
elements in their atmosphere?
So far, no star has been observed that has no iron or
other heavy elements in its atmosphere. But the current
theories of Big Bang cosmology predict there must have
been some stars which were composed almost only of
hydrogen and helium. These are known as population III
stars.
A few minutes after the Big Bang there was a period of
primordial nucleosynthesis: the creation of elements
from the sea of protons and neutrons that filled the
Universe. This resulted in the production of deuterium
(a hydrogen isotope of one proton and one neutron),
helium-3 and -4, and lithium-6 and -7. It is from these
elements that the first stars would have formed.
These population III stars are predicted to be much more
massive than stars that exist today. Model simulations
predict that they would have been over 100 times the
mass of the Sun. Due to their massive size they would
have lived for only a couple of million years before
exploding in a supernova. For stars in the mass range
of 130 solar masses, this would result in a supernova
for which is predicted a large proportion of the heavy
elements created would take the form of iron.
Population III stars have not been observed yet. This is
due to the extremely short lifetimes and existing only in
the very early Universe. The most metal-poor stars that
are observed today are found in groups of stars called
globular clusters. These are spherical conglomerations
of between 10,000 and a million stars that form halos
around galaxies.
One of the stars with the lowest known amount of heavy
elements is HE0107-5240, with about 1/200,000 of the
iron content of the Sun (Lau, Stancliffe and Tout).
What about the other extreme? That is, stars that have an
Figure 6: The spiral galaxy NGC 2770 and its two
supernovae. The bright star at the edge of the galaxy in
the top right is SN 2008D, while the star left of the centre is
SN 2007uy.
Ref: http://www.eso.org/public/images/eso0823a/
equivalently high (200,000 times) abundances of iron?
No such stars are known, though it is theorized by some
that the surfaces of neutron stars could consist of iron.
Neutron stars are the remnant of supernova. They consist
of neutrons supported by the Pauli Exclusion Principle.
But their surface regions are thought to be composed
of atomic nuclei, with the possibility of it being iron left
from the core of the massive star that underwent the
supernova.
Conclusion
Through the interplay of the strong nuclear force and
electromagnetic force, iron finds itself with one of the
highest binding energies per nucleon. In addition, iron56 can be created by the addition of alpha particles in
the interior of stars allowing it to be built up in the core
of massive stars. It is here that iron acts as one of the final
fusion products. With no more energy available, the star’s
core collapses, causing a supernova.
For further information contact:
[email protected]
Bibliography
Fewell, M. P. (1995). “The atomic nuclide with the highest mean binding energy.”
American Journal of Physics, 63.7, 653-658.
Fleurot, F. (2010). Evolution of Massive Stars. 20 April 2010, http://nu.phys.
laurentian.ca/~fleurot/evolution/
Halliday, D., Resnick, R., & Walker, J. (2005). Fundamentals of Physics. 7th Edition.
Wiley.
Lau, Herbert H.B., Stancliffe, R. J., & Tout, C. A. (2007). “Carbon-rich extremely metal
poor stars: signatures of Population III asymptotic giant branch stars in binary
systems.” Monthly Notices of the Royal Astronomical Society, 378.2, 563-568.
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