Download Scientific Background Paper: Iron`s Place and Role

Scientific Background Paper: Iron’s Place and Role in the Universe
The formation of heavy elements in the Universe is a by-product of the different processes that fuel stars
during the various stages of their evolution and ultimately their death. The process of forming nuclei is
known as nucleosythesis and when this takes place in stars it is referred to as stellar nucleosynthesis. The
relative importance of each process in the chain of events that describes the phases of stellar evolution
depends mainly on the initial conditions (mass and composition) present at the birth of the star. These
initial conditions in turn depend on the local star formation history and their study remains an active field
of research in stellar astrophysics.
In order to understand the special place that Iron holds in stellar evolution, it is important to consider the
complete chain of events that lead up to the formation of Iron nuclei. Iron plays a key role in the normal
evolution of stars and is further produced in vast quantities in the subsequent explosive event at the end of
the process of stellar evolution, which gives rise to the metals we see in the Universe today, including
those found on earth. It is no exaggeration to say that the process of stellar evolution is arguably the most
important chain of events in the Universe that we know of. It is responsible for making the Universe the
way it appears today and our understanding of it has improved our knowledge of the evolution of the
universe immensely.
Big Bang nucleosynthesis:
All of the light elements that would later be used as fuel for the first generation of stars were formed in
the first few minutes after the Big Bang. Once the Universe had cooled sufficiently to allow the formation
of Protons and Neutrons (about 100 seconds after the Big Bang), the first stage of nucleosynthesis could
take place. This ‘Big Bang Nucleosynthesis’ took the form of proton and neutron capture forming
successively heavier elements. However, this process ended almost as quickly as it had begun (about 1000
seconds after the Big Bang) as the temperature of the early Universe fell and the nucleons no longer had
sufficient kinetic energy to overcome the Coulomb barrier. This lead to a drastic reduction in the reaction
rates for the relevant processes and nucleosynthesis effectively stopped. The early work on Big Bang
Nucleosynthesis was highlighted in the famous ‘ ’ paper (Alpher, Bethe and Gamow 1948).
This process failed to form heavy elements due to the lack of a stable isotope with mass numbers 5 or 8.
This led to the primordial elemental abundances of Hydrogen, Deuterium, Helium-3, Helium-4 and
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Lithium-7 seen today and provided the fuel for the second stage of element formation, the process of
stellar nucleosynthesis.
Stellar formation and birth:
The stages of stellar evolution are determined by the mechanism used to limit the collapse of a star by
producing an outward force to overcome the inward pull of gravity. The lifecycle of stars begins with the
in-fall of a ball of gas and dust and their death is ultimately the failure to find a mechanism to arrest this
in-fall. The details of the mechanisms used and the various stages of the evolution during the lifecycle are
determined by the mass and composition of the initial gas cloud from which they form.
In general, stars are formed when a cloud of gas collapses under the force of gravity. The process is
unstable and once begun the collapse accelerates until the cloud occupies a small enough volume that the
pressure due to collisions between particles begins to slow its inward collapse. The star comes into
thermodynamic equilibrium where the system can be described in terms of it bulk statistical properties.
The Virial theorem for a non-degenerate gas tells us that half of the gravitational energy converted from
the contraction of the gas is converted into the form of kinetic energy of the particles and the remainder is
consumed by the production of radiation. Initially, the photons ionize the contracting gas allowing the
temperature to rise in the core. Once the gas is sufficiently ionized, the star enters hydrostatic equilibrium
where the outward pressure of the gas balances the inward gravitational force. At this stage the star is
approaching the Main Sequence and continues to collapse slowly until it reaches it and thermonuclear
reactions can occur.
Once the star has contracted sufficiently to reach this stage the star becomes relatively static with the
thermonuclear processes slowly changing the chemical composition of the star. The change in chemical
composition begins with the conversion of Hydrogen into Helium. Stars on the Main Sequence are
burning hydrogen in their cores and evolve slowly as they build up a mass of Helium, slowly depleting the
Hydrogen fuel.
In order for two nuclei to fuse they must overcome the repulsive force of the electrostatic (or Coulomb)
repulsion between the nuclei. This repulsive force acts at large distances and the nuclei must overcome it
in order to get sufficiently close so that the attractive nuclear force dominates and the nucleons become
bound. The Coulomb barrier is ~1MeV which is far greater than the kinetic energy of the nuclei in the
core of the star. However, the process is possible due to barrier penetration, a quantum mechanical effect
that means there is a finite probability that the nuclei can exist at a separation where the nuclear force
dominates. Therefore, in an ensemble of nuclei a finite number of nuclei will fuse. The probability that
two nuclei will fuse is strongly dependent on their kinetic energy and hence the temperature of the gas.
This higher the kinetic energy of the nuclei, the lower the potential barrier they must tunnel through and
hence the higher the probability of success.
The height of the Coulomb barrier is a function of the charge of the nuclei, which is in turn a function of
the mass of the nucleus (i.e. the number of nucleons). Hence, the temperature needed to fuse heavier
nuclei increases steadily, which limits core burning to light nuclei in all but the most massive stars.
Stellar Nucleosynthesis:
The light elements are recycled into heavier elements by thermonuclear reactions. The exact process was
first put forward in the groundbreaking paper of (Burbidge, et al. 1957). The process depends on the
conditions in the star and change during the lifetime of the star. The star begins life with Hydrogen as
the most abundant element and slowly builds core of heavier nuclei, which will become the fuel for the
subsequent stages. This building of heavier nuclei continues throughout the lifetime of the star, increasing
the amount of heavier elements in the core until a limit is reached with the iron-group of nuclei, where
the process becomes energetically unfavorable. The detailed rates and total mass of each type of nuclei
vary enormously depending on the initial mass and composition of the star. They are constantly changing
depending on the conditions to be found in the various layers of the star. In the following sections I will
describe some of the processes occurring during the lifetime of a star, focusing on those that are relevant
to the formation of the iron-group elements. More detailed explanations of the various processes can be
found in, for example, (Clayton 1968).
The first stages of nucleosynthesis are concerned with the fusion of Hydrogen into Helium and the
formation of a Helium core. This stage is known as the Main Sequence for stars and the duration of this
phase depends on the initial mass and composition of the star (heavier, more ‘metal’ rich stars evolve
faster).
The proton-proton chain
The first important process for the formation of Helium takes place via the proton-proton (or p-p) chain
and was first explained in depth by (Bethe 1939), leading to him being awarded the Nobel Prize for
Physics(Salpeter 1952). Instead of the simple fusion of multiple Hydrogen nuclei via capture, which
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would have an incredibly low reaction rate, the reaction takes the form of a change of processes forming a
reaction chain. The first reaction is the conversion of two hydrogen nuclei (i.e. protons) into a single
Deuterium nucleus (proton + neutron), which also liberates a positron and a neutrino. The resulting
Deuterium nucleus then captures a third Hydrogen nucleus, forming a Helium-3 nucleus and liberating
an energetic photon whose energy is passed to the surrounding particles. In the p-p I chain, two Helium3 nuclei subsequently fuse to form a Helium-4 nucleus, which in turn liberate 2 protons. The p-p II and
p-p III chains rely on a sufficient amount of Helium having already been formed and involve the Helium3 nucleus fusing with a Helium-4 nucleus to form a Beryllium-7 nucleus. If the Beryllium-7 nucleus can
capture an electron it will become Lithium-7, which in turn fuses with a Hydrogen nucleus to become
two Helium-4 nuclei. This process is the p-p II chain. In the p-p III chain, the Beryllium-7 nucleus
captures a proton and becomes Boron-8, which decays into Beryllium-8 and subsequently into two
Helium-4 nuclei.
The relative importance of the chains depends on the temperature, density and abundances of the nuclei
involved. The rate of each chain changes as the abundance and temperature of the stellar interior change
but the final product for each chain is the same and the outcome of this stage of stellar nucleosynthesis is
the same, the formation of Helium-4 nuclei. The entire chain has used 4 Hydrogen nuclei (protons) and
created a single Helium-4 nucleus (also called an alpha-particle). The reaction rate of the any process is
limited by the slowest reaction. In the case of the p-p chain this reaction is the formation of Deuterium,
the timescale for which is almost 1010 years and is only feasible as an energy source due to the large
number of nuclei present in the core of a star.
The Carbon-Nitrogen-Oxygen (CNO) cycle:
In stars where the initial composition is not purely Hydrogen, the presence of Carbon, Nitrogen and
Oxygen nuclei can induce a chain of reactions that again convert Hydrogen nuclei into Helium nuclei,
liberating enough energy to arrest the contraction of the star under self-gravitation in the process. The
CNO nuclei act similarly to a catalyst in a chemical reaction and are not generated nor destroyed in the
process. This then relies on the prior composition of CNO nuclei to proceed. If there is sufficient CNO
present, there are two processes that can proceed. The first begins with the capture of a Hydrogen nucleus
by a Carbon-12 nucleus, leading to the formation of a Nitrogen-13 nucleus, which subsequently beta
decays to a Carbon-13 nucleus. The Carbon-13 captures a further Hydrogen nucleus to become
Nitrogen-14, which again captures a Hydrogen nucleus to become Oxygen-15 before beta decaying to
Nirogen-15. Finally, the Nitrogen-15 nucleus captures a Hydrogen nucleus and becomes a Carbon-12
and a Helium-4 nucleus. Occasionally, Oxygen-16 is produced during the capture of Hydrogen on
Nitrogen-15. In this case, further proton-capture leads to Fluorine-17 production, which beta decays to
Oxygen-17 before capturing another Hydrogen nucleus and returning to Nitrogen-14 and a Helium-4
nucleus. These processes each convert four Hydrogen nuclei into a single Helium-4 nucleus and are
important at higher temperatures than the p-p cycle. The processes involved in the CNO cycle have
stronger temperature dependences and will begin to dominate when the temperature is sufficiently high.
The p-p and CNO cycles are the characteristic processes taking place during the main sequence lifetime
of a star. In addition to providing the energy to support the star, they produce the fuel (Helium-4 nuclei)
for all further stages of stellar nucleosynthesis. As the Hydrogen is converted into Helium, the density of
Helium-4 increases and the core burning of heavier nuclei can start. This process occurs in stars with a
mass greater than about 0.5 solar-masses and begins the next stage of stellar evolution and the next stage
of stellar nucleosynthesis.
Helium burning:
The reason that a star builds a core of Helium-4 and does not simply burn Helium-4 at a steady rate
matching its production is the lack of a stable nucleus with atomic number of 8. The Beryllium-8 isotope
created by the fusion of two Helium-4 nuclei has a lifetime of 2.6 x 10-16 seconds. However, once a
sufficient amount of Helium-4 builds up and the temperature is sufficiently high, the Beryllium-8 can
capture another Helium-4 nucleus leading to the formation of a Carbon-12 nucleus. This is known as the
triple-alpha process (Salpeter 1952) and it is the increased cross-section of the capture by an intermediate
Beryllium-8 nucleus compared to that of a single interaction of three Helium-4 nuclei that makes the
reaction feasible as a source of energy. In addition to the formation of Carbon-12, which is thought to be
responsible for high abundance of Carbon-12 in the universe, the Carbon-12 can also subsequently
capture another Helium-4 nucleus and form Oxygen-16, the next most abundant nuclei in the Universe.
In a manner similar to the generation of Helium-4 nuclei from Hydrogen nuclei, the ashes of Helium-4
burning become the fuel for the next stage in stellar evolution for stars of sufficient mass (at least
approximately eight solar masses).
Carbon and Oxygen burning:
Beginning with Carbon burning the process of building up heavier elements becomes more complex and
less predictable. The fusion of two Carbon-12 nuclei occurs in stars with a mass of more than around 0.7
solar masses and leads to the production of an excited nucleus that decays to a more stable configuration
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leading to the production of Magnesium-24, Magnesium-23, Sodium-23, Neon-20 and further Oxygen16. The decay products are also quickly captured leading to further reactions. Each reaction has its own
rate that is temperature dependent and the situation in the stellar core changes rapidly. (Clayton 1968)
gives a review of all of the reaction processes of Carbon and Oxygen burning and rates and how these
rates are determined.
In the case of Oxygen-16 burning, which occurs at higher temperatures than the Carbon-12 burning, the
product is again in an excited state and quickly decays to different, heavier products, including Sodium32, Sodium-31, Phosphorus-31, Silicon-28 and Magnesium-24. The most important of these is Silicon
and this plays an important role in the final stages of stellar evolution.
Silicon burning:
In the final stages of stellar nucleosynthesis, the fusing of two Silicon-28 nuclei into an Iron group
nucleus is energetically unfavorable and instead photodisintegration and Helium-4 capture become the
dominant processes. This process occurs in massive stars and lasts only a few weeks compared to the
millions of years of the evolution up to this point. The process begins when a high-energy gamma-ray
strips a light nucleus from a heavy target nucleus (e.g. Sulphur-32). The light nucleus could be a proton, a
neutron or an alpha particle (a Helium-4 nucleus). These light nuclei then rapidly combine with another
target nucleus. This could be the original nucleus, or one of the other nuclei of a different species present
in the core. This rapid exchange of light particles via photodisintegration leads to a wide variety of nuclei
being formed and a constant exchange. The process reaches a semi-stable state with the balance of target
nuclei and light nuclei being generated and destroyed, leading to what is known as nuclear statistical
equilibrium. The ultimate abundances of the intermediate mass nuclei depend on effects such as
convective mixing in the later stages of stellar evolution, which are difficult to model. In contrast to
subsequent stages of core burning, the intermediate mass elements made in these nuclear reactions remain
relatively unchanged during the subsequent supernova explosion and reflect the ejected abundances
(Woosley and Weaver 1995).
The most important process in building large nuclei is the capture of an alpha particle and the subsequent
emission of a photon, e.g. Silicon-28 + Helium-4, which produces Sulphur-32. While the equilibrium
state is a process of generation and disintegration, there is a tendency for the process to increase the
number of heavy element nuclei. However, instead of continuing indefinitely, the process stalls when the
binding energy per nucleon reaches a maximum (i.e. where fusion with lighter elements leads to the net
release of energy). This occurs in the chain of alpha particle captures with Iron-56 (i.e. the production of
Zinc-60 is exothermic) and, for this reason, the process of Silicon burning continues up to the formation
of the iron group nuclei (Iron, Cobalt and Nickel). The dominant nuclei present at the end of this process
are in fact Iron-54 and Nickel-56. However, the Nickel-56 nucleus is unstable and beta decays to Cobalt56 and then decays again to Iron-56, with a half-life of 6 days and 77 days respectively. This leads to a
build up of a core of the two most common isotopes of Iron, Fe-54 and Fe-56.
Evolution beyond Silicon burning:
Once the star has built up a core of iron group nuclei, there is no longer a fuel source capable of
supporting the energy generation needed to support the star and the star is destined to become a corecollapse supernova (Woosley and Janka 2005). Alpha particle captures beyond the iron group take energy
from the core and photodisintegration of iron group nuclei into alpha particles would require up to
100MeV of energy per iron nucleus. The star begins to contract and the core temperature rises. This leads
to further photodisintegrations, removing photons and also to the capture of electrons by heavy nuclei,
reducing the pressure and allowing the core to contract further. As the photodisintegration and electron
capture processes continue, it becomes a runaway process and the star begins to contract more and more
rapidly until it is almost in a free-fall collapse.
At this stage a large number of neutrons are being produced and, for massive stars, the process is only
halted by the neutron degeneracy pressure, which leads to the formation of a neutron star at the center of
the star. This sudden halting of the in-fall leads to the formation of an outward travelling shockwave,
which deposits energy into the unburned shells of material formed during the previous stages of evolution.
The outer layers of the star are still processing lighter nuclei into iron group elements, which are
subsequently ejected into the interstellar medium by the explosion. This type of explosion is known as a
Type-II supernova and is the fate of massive stars, with initial masses above about 20 times the mass of
the sun. The outward travelling shockwave deposits energy in the material surrounding the core raising
the temperature and leading to further ‘burning’ of the nuclear fuel. This is followed by a burst of
neutrinos from the core, which deposits further energy. The energy deposition is sufficient to fuel further
Silicon burning, leading to the generation of Nickel-56 in large quantities, which is ejected into the
interstellar medium.
The exact mass of ejected material and its composition is difficult to predict and involves detailed time
dependent calculations of nuclear reaction networks (Woosley and Weaver 1995). The radioactive decay
of Nickel-56 into Cobalt-56 and subsequently into Iron-56 generates photons which are the dominant
source of the observed lightcurves of supernovae following the initial explosion and explain the
exponential decay of the flux of photons from the supernova, as the natural radioactive decay of the nuclei
forms the stable Iron-56 nuclei. The ejected mass of Nickel-56 (and hence Iron-56) is 1029 kg per
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explosion for an initial 25 solar mass star. The rate of core-collapse supernovae is approximately 3 per
century in our galaxy.
Recently, the role of Type Ia supernova in stellar nucleosynthesis has begun to be explored. Type Ia
supernovae are the explosion of a Carbon+Oxygen white dwarf that has exceeded the Chandrasekhar
limit, the maximum mass for a C/O white dwarf ( 1.1 solar masses). The evolutionary track for these
systems is the topic of recent debate, but the result in both cases is a detonation of ~1 solar mass of
Carbon and Oxygen. The exact details of the nucleosynthesis remain poorly understood, but involve the
same processes as occur in the explosive nucleosynthesis of core-collapse supernovae (Nuclear statistical
equilibrium, Si, Oxygen and Carbon burning). Whilst the individual yield is high (up to 0.6 solar masses
of Ni-56 per system), it is thought they could be responsible for a similar amount of Ni-56 as produced in
Type II events when their lower rate is taken into account.
A notable example was the recent supernova SN2011fe in the nearby galaxy M101, which was visible
with small telescopes from sites in the northern hemisphere in 2011 (Wikipedia 2011).
Heavy element nucleosynthesis:
The formation of nuclei beyond the iron group (the so-called ‘Heavy Elements’) is only possible in the
presence of free neutrons and is the result of neutron capture onto iron group nuclei. The neutrons are
captured building up heavier and heavier nuclei(Burbidge, et al. 1957). Eventually, these heavy nuclei will
be unstable to beta decay and the nucleus will either decay or capture a subsequent neutron depending on
the relative rates of the two processes. These processes for building heavier nuclei are called the slow
process (or s-process) and the rapid process (or r-process) depending on whether the neutron capture
timescale is slower or more rapid than the beta decay timescale respectively. These processes are
responsible for forming all of the elements heavier than the iron group, but their relatively slow reaction
rates make them only important during the explosive nucleosynthesis of a core-collapse (Type-II)
supernova. This in turn explains the relatively low abundances of such elements.
Iron abundance in the Universe:
The chemical enrichment of the Universe is a direct result of the process of stellar nucleosynthesis. The
enrichment is dominated by the input from supernovae explosions in all but the most massive stars, where
losses from stellar winds must also be taken into account. However, such wind driven losses contribute
little to the overall abundance. The result of both processes is however the same; the composition of
galaxies is gradually changing, with a steady increase in the amount of heavy elements. As the fuel for
newly formed stars is composed of the ashes from previous supernova explosions, the initial composition
of stars is also evolving with time. This has led to the division of stars into three groups, the so-called
population I, II and III stars. Population I stars are the youngest stars and have an increased amount of
elements heavier than Lithium. Populations II stars are those for which little enrichment has taken place
and Population III stars (which have yet to be observed) are postulated to be the first stars to form after
the big-bang and would therefore have primordial abundances. Iron is thought to be the sixth most
abundant element in the Universe with ~0.1% by mass, which is also consistent with the solar system
elemental abundances.
The connection to planet formation:
There is growing evidence that planets preferentially form around stars with high metallicity (Johnson
and H. 2012). The first generation of stars (Population III stars) would likely not have had planets form
and it was only following the first supernovae eruptions that the Universe became enriched enough for
planets to form. In the popular core accretion model of planet formation the dust act as seeds for building
planetesimals, which further combine and accrete further dust grains to form planets and eventually the
planetary systems(Blum and Wurm 2008). This process if though to be the mechanism for the formation
of planets in our solar system and the growing number observed beyond our solar system by missions such
as NASA’s ground-breaking Kepler mission(NASA Kepler mission website 2012). The grains are
particles of the elements created during nucleosythesis and are rich in iron, reflecting the yield of the
supernova explosion.
The connection to planet habitability:
As the number of planets discovered orbiting around other stars has steadily grown since the discovery of
the first extra-solar planet in orbiting 51 Pegasi (Mayor and Queloz 1995) to its current level of 763 (The
Extrasolar Planets Encyclopaedia 2012), the focus of the search has shifted from detecting planets to
detecting planets capable of harboring life. This has been advanced greatly by NASA’s Kepler satellite,
whose mission is to detect potentially life-supporting planets around other stars (NASA Kepler mission
website 2012). To date, Kepler has detected 61 confirmed planets and has over 2000 further candidates of
which more than 900 are thought to have a radius less than twice that of the Earth. Due to the limitations
of current technology, it is only possible to derive basic information on the planets and planet candidates
and will remain so for some time to come. In order to compare the potential for harboring life between
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the ever-increasing number of planets, a number of concepts and metrics have been introduced. Whilst
there are many interpretations of these parameters, reflecting amongst other things, the different
interpretations of the form of life being considered, the most commonly accepted versions have been
collated by the Planetary Habitability Laboratory of the University of Puerto Rico, Arecibo. These are;
The Habitable (Goldilocks) zone:
(Kasting, Whitmire and Reynolds 1993) states that “For Earth-like planets orbiting main sequence stars,
the HZ is determined by water loss on the inner edge and by CO2 condensation, leading to runaway
glaciation, on the outer edge”. This has become known colloquially as the ‘Goldilocks zone’, where it is
‘not too hot’ on the inner edge and ‘not too cold’ on the outer edge.
Earth Similarity Index (ESI):
The Earth Similarity Index is an attempt to parameterize the properties that combine to make a planet
more similar or not to Earth (Schulze-Makuch 2011). It takes into account a range of criteria for
characterizing how similar a planet is to earth and therefore how capable of support similar forms of life
to those found on earth. The variables used are planetary radius, density, escape velocity and surface
temperature. The variables are combined with weights in order to attempt to take into account the
importance of each variable to the overall similarity. Initial studies have taken place to use data from the
currently known and candidate planets to form a consensus on the frequency of ‘Earth Similarity’
(Schulze-Makuch 2011).
Planetary Habitability Index (PHI):
(Schulze-Makuch 2011) have gone further and defined the PHI to assess ‘the possibility that life in some
form could be present on another world’. The show that life requires ‘a stable and protected substrate,
energy, polymeric chemistry, and a liquid medium’. These include factors such as how rocky a planet is,
the presence of an atmosphere, the presence of a magnetosphere to shield the planet surface, the existence
of a suitable energy source, the presence of Carbon, Nitorgen, Sulphur and Phosphorus and the presence
of liquids (most importantly water) in the atmosphere, on the surface, or below the surface.
Habitable Zone Composition (HZC):
The HZC is a measure of the bulk composition of a planet. The scale has a value -1 when the planet has
a mean density equal to that of iron and +1 when it has a mean density equal to that of water. For
reference, Earth has a HSC of -0.31, as would be expected of a rocky planet with an iron core.
Bibliography
Alpher, R., H. Bethe, and G. Gamow. “The origin of chemical elements.” Physical Review, no. 73 (1948):
803.
Bethe, H. “Energy Production in Stars.” Physical Review 55 (1939): 434.
Blum, J., and G. Wurm. “The Growth Mechanisms of Macroscopic Bodies in Protoplanetary Disks.”
Annual Reviews of Astronomy and Astrophysics, 2008.
Burbidge, M., G Burbidge, W Fowler, and F Hoyle. “Synthesis of the elements in stars.” Reviews of
modern physics 29 (1957): 547.
Clayton, D. Principles of Stellar Evolution and Nucleosynthesis. Universioty of Chicago Press, 1968.
Johnson, J., and Li H. “The First Planets: the Critical Metallicity for Planet Formation.” Astrophysical
Journal, 2012.
Kasting, J., D. Whitmire, and R. Reynolds. “Habitable Zones around Main Sequence Stars.” Icarus 101
(1993): 108.
Mayor, M., and D Queloz. “A jupiter-mass companion to a solar-type star.” Nature 6555 (1995): 355.
NASA Kepler mission website. Kepler Homepage. 2012. www.nasa.gov/kepler/.
Salpeter, E. “Nuclear Reactions in Stars Without Hydrogen.” Astrophysical Journal 115 (1952): 326.
Schulze-Makuch, D. “A Two-Tiered Approach to Assessing the Habitability of Exoplanets.” Astrobiology
11 (2011).
The Extrasolar Planets Encyclopaedia. 2012 йил May. http://exoplanet.eu/catalog.php.
Wikipedia. 2011. http://en.wikipedia.org/wiki/SN_2011fe.
Woosley, S, and T. Weaver. "The evolution and explosion of massive stars. II. Explosive hydrodynamics
and nucleosynthesis." The Astrophysical Journal Supplement Series 101 (1995): 181.
Woosley, S., and T. Janka. “The physics of core collapse supernovae.” Nature 1 (2005): 147.
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