Carbon-burning process

The carbon-burning process or carbon fusion is a set of nuclear fusion reactions that take place in the cores of massive stars (at least 8 at birth) that combines carbon into other elements. It requires high temperatures (> 5×108 K or 50 keV) and densities (> 3×109 kg/m3).[1]

These figures for temperature and density are only a guide. More massive stars burn their nuclear fuel more quickly, since they have to offset greater gravitational forces to stay in (approximate) hydrostatic equilibrium. That generally means higher temperatures, although lower densities, than for less massive stars.[2] To get the right figures for a particular mass, and a particular stage of evolution, it is necessary to use a numerical stellar model computed with computer algorithms.[3] Such models are continually being refined based on nuclear physics experiments (which measure nuclear reaction rates) and astronomical observations (which include direct observation of mass loss, detection of nuclear products from spectrum observations after convection zones develop from the surface to fusion-burning regions – known as dredge-up events – and so bring nuclear products to the surface, and many other observations relevant to models).[4]

Fusion reactions

The principal reactions are:[5]

12
6
C
 
12
6
C
 
→  20
10
Ne
 
4
2
He
 
4.617 MeV
12
6
C
 
12
6
C
 
→  23
11
Na
 
1
1
H
 
2.241 MeV
12
6
C
 
12
6
C
 
→  23
12
Mg
 
1n  −  2.599 MeV
Alternatively:
12
6
C
 
12
6
C
 
→  24
12
Mg
 

γ
 
13.933 MeV
12
6
C
 
12
6
C
 
→  16
8
O
 
4
2
He
 
−   0.113 MeV

Reaction products

This sequence of reactions can be understood by thinking of the two interacting carbon nuclei as coming together to form an excited state of the Mg-24 nucleus, which then decays in one of the five ways listed above.[6] The first two reactions are strongly exothermic, as indicated by the large positive energies released, and are the most frequent results of the interaction. The third reaction is strongly endothermic, as indicated by the large negative energy indicating that energy is absorbed rather than emitted. This makes it much less likely, yet still possible in the high-energy environment of carbon burning.[5] But the production of a few neutrons by this reaction is important, since these neutrons can combine with heavy nuclei, present in tiny amounts in most stars, to form even heavier isotopes in the s-process.[7]

The fourth reaction might be expected to be the most common from its large energy release, but in fact it is extremely improbable because it proceeds via electromagnetic interaction,[5] as it produces a gamma ray photon, rather than utilising the strong force between nucleons as do the first two reactions. Nucleons look a lot bigger to each other than they do to photons of this energy. However, the Mg-24 produced in this reaction is the only magnesium left in the core when the carbon-burning process ends, as Mg-23 is radioactive.

The last reaction is also very unlikely since it involves three reaction products,[5] as well as being endothermic — think of the reaction proceeding in reverse, it would require the three products all to converge at the same time, which is less likely than two-body interactions.

The protons produced by the second reaction can take part in the proton-proton chain reaction, or the CNO cycle, but they can also be captured by Na-23 to form Ne-20 plus a He-4 nucleus.[5] In fact, a significant fraction of the Na-23 produced by the second reaction gets used up this way.[6] In stars between 9 and 11 solar masses, the oxygen (O-16) already produced by helium fusion in the previous stage of stellar evolution manages to survive the carbon-burning process pretty well, despite some of it being used up by capturing He-4 nuclei.[1][8] So the end result of carbon burning is a mixture mainly of oxygen, neon, sodium and magnesium.[3][5]

The fact that the mass-energy sum of the two carbon nuclei is similar to that of an excited state of the magnesium nucleus is known as 'resonance'. Without this resonance, carbon burning would only occur at temperatures one hundred times higher. The experimental and theoretical investigation of such resonances is still a subject of research.[9] A similar resonance increases the probability of the triple-alpha process, which is responsible for the original production of carbon.

Neutrino losses

Neutrino losses start to become a major factor in the fusion processes in stars at the temperatures and densities of carbon burning. Though the main reactions don't involve neutrinos, the side reactions such as the proton-proton chain reaction do. But the main source of neutrinos at these high temperatures involves a process in quantum theory known as pair production. A high energy gamma ray which has a greater energy than the rest mass of two electrons (mass-energy equivalence) can interact with electromagnetic fields of the atomic nuclei in the star, and become a particle and anti-particle pair of an electron and positron.

Normally, the positron quickly annihilates with another electron, producing two photons, and this process can be safely ignored at lower temperatures. But around 1 in 1019 pair productions[2] end with a weak interaction of the electron and positron, which replaces them with a neutrino and anti-neutrino pair. Since they move at virtually the speed of light and interact very weakly with matter, these neutrino particles usually escape the star without interacting, carrying away their mass-energy. This energy loss is comparable to the energy output from the carbon fusion.

Neutrino losses, by this and similar processes, play an increasingly important part in the evolution of the most massive stars. They force the star to burn its fuel at a higher temperature to offset them.[2] Fusion processes are very sensitive to temperature so the star can produce more energy to retain hydrostatic equilibrium, at the cost of burning through successive nuclear fuels ever more rapidly. Fusion produces less energy per unit mass as the fuel nuclei get heavier, and the core of the star contracts and heats up when switching from one fuel to the next, so both these processes also significantly reduce the lifetime of each successive fusion-burning fuel.

Up to the helium burning stage the neutrino losses are negligible. But from the carbon burning stage onwards, the reduction in stellar lifetime due to energy lost in the form of neutrinos roughly matches the increased energy production due to fuel change and core contraction. In successive fuel changes in the most massive stars, the reduction in lifetime is dominated by the neutrino losses. For example, a star of 25 solar masses burns hydrogen in the core for 107 years, helium for 106 years and carbon for only 103 years.[10]

Stellar evolution

During helium fusion, stars build up an inert core rich in carbon and oxygen. The inert core eventually reaches sufficient mass to collapse due to gravitation, whilst the helium burning moves gradually outward. This decrease in the inert core volume raises the temperature to the carbon ignition temperature. This will raise the temperature around the core and allow helium to burn in a shell around the core.[11] Outside this is another shell burning hydrogen. The resulting carbon burning provides energy from the core to restore the star's mechanical equilibrium. However, the balance is only short-lived; in a star of 25 solar masses, the process will use up most of the carbon in the core in only 600 years. The duration of this process varies significantly depending on the mass of the star.[12]

Stars with below 8–9 Solar masses never reach high enough core temperature to burn carbon, instead ending their lives as carbon-oxygen white dwarfs after shell helium flashes gently expel the outer envelope in a planetary nebula.[3][13]

In stars with masses between 8 and 11 solar masses, the carbon-oxygen core is under degenerate conditions and carbon ignition takes place in a carbon flash, that lasts just milliseconds and disrupts the stellar core.[14] In the late stages of this nuclear burning they develop a massive stellar wind, which quickly ejects the outer envelope in a planetary nebula leaving behind an O-Ne-Na-Mg white dwarf core of about 1.1 solar masses.[3] The core never reaches high enough temperature for further fusion burning of heavier elements than carbon.[13]

Stars with more than 11 solar masses start carbon burning in a non-degenerate core,[14] and after carbon exhaustion proceed with the neon-burning process once contraction of the inert (O, Ne, Na, Mg) core raises the temperature sufficiently.[13]

See also

References

  1. ^ a b Ryan, Sean G.; Norton, Andrew J. (2010). Stellar Evolution and Nucleosynthesis. Cambridge University Press. p. 135. ISBN 978-0-521-13320-3.
  2. ^ a b c Clayton, Donald (1983). Principles of Stellar Evolution and Nucleosynthesis. University of Chicago Press. ISBN 978-0-226-10953-4.
  3. ^ a b c d Siess L. (2007). "Evolution of massive AGB stars. I. Carbon burning phase". Astronomy and Astrophysics. 476 (2): 893–909. Bibcode:2006A&A...448..717S. doi:10.1051/0004-6361:20053043.
  4. ^ Hernandez, G.; et al. (Dec 2006). "Rubidium-Rich Asymptotic Giant Branch Stars". Science. 314 (5806): 1751–1754. arXiv:astro-ph/0611319. Bibcode:2006Sci...314.1751G. doi:10.1126/science.1133706. PMID 17095658.
  5. ^ a b c d e f Camiel, W. H.; de Loore; C. Doom (1992). Camiel W. H. de Loore (ed.). Structure and evolution of single and binary stars. Astrophysics and Space Science Library. Springer. pp. 95–97. ISBN 978-0-7923-1768-5.
  6. ^ a b Rose,, William K. (1998). Advanced Stellar Astrophysics. Cambridge University Press. pp. 227–229. ISBN 978-0-521-58833-1.
  7. ^ Rose (1998), pp. 229–234
  8. ^ Camiel (1992), pp.97–98
  9. ^ Strandberg, E.; et al. (May 2008). "24Mg(α,γ)28Si resonance parameters at low α-particle energies". Physical Review C. 77 (5): 055801. Bibcode:2008PhRvC..77e5801S. doi:10.1103/PhysRevC.77.055801.
  10. ^ Woosley, S.; Janka, H.-T. (2006-01-12). "The Physics of Core-Collapse Supernovae". Nature Physics. 1 (3): 147–154. arXiv:astro-ph/0601261. Bibcode:2005NatPh...1..147W. CiteSeerX 10.1.1.336.2176. doi:10.1038/nphys172.
  11. ^ Ostlie, Dale A. and Carrol, Bradley W., An introduction to Modern Stellar Astrophysics, Addison-Wesley (2007)
  12. ^ Anderson, Scott R., Open Course: Astronomy: Lecture 19: Death of High-Mass Stars, GEM (2001)
  13. ^ a b c Ryan (2010), pp.147–148
  14. ^ a b The Carbon Flash
B2FH paper

The B2FH paper, named after the initials of the authors of the paper, Margaret Burbidge, Geoffrey Burbidge, William A. Fowler, and Fred Hoyle, is a landmark paper on the origin of the chemical elements published in Reviews of Modern Physics in 1957. The title of that paper is "Synthesis of the Elements in Stars", but as that paper grew in influence, it came to be referred to only as "B2FH". The B2FH paper spread stellar nucleosynthesis theory widely in the scientific community, especially among astronomers who saw everyday relevance to their quest, at a time when it was appreciated by only a handful of experts in nuclear physics. But it did not create the theory of stellar nucleosynthesis as much as bring it vividly to life.

The paper comprehensively outlined and analyzed several key processes that are responsible for the nucleosynthesis of the elements heavier than iron and their relative abundance by the capture within stars of free neutrons. It advanced much less the understanding of the synthesis of the very abundant elements from silicon to nickel. A puzzle about that is that despite Hoyle's coauthorship of B2FH and being its chief conceptual architect, the paper did not include the carbon-burning process, the oxygen-burning process and the silicon-burning process, each of which contributes massively to the growth of stellar metallicity from magnesium to nickel in the interstellar gas. The supernova nucleosynthesis that achieves that had been published by Hoyle in 1954. Donald D. Clayton has attributed the severe undercitations of Hoyle's 1954 paper relative to the voluminous citations of B2FH to several factors: the advanced difficulty of digesting Hoyle's 1954 paper even for his B2FH coauthors, as it proved to be for the world of astronomy generally; to Hoyle's having described its key equation only in words rather than writing it prominently in his paper; and finally to a lack of careful review by Hoyle himself of the B2FH draft written by two junior coauthors who had themselves not adequately digested Hoyle's paper.

Index of physics articles (C)

The index of physics articles is split into multiple pages due to its size.

To navigate by individual letter use the table of contents below.

Neon

Neon is a chemical element with the symbol Ne and atomic number 10. It is a noble gas. Neon is a colorless, odorless, inert monatomic gas under standard conditions, with about two-thirds the density of air. It was discovered (along with krypton and xenon) in 1898 as one of the three residual rare inert elements remaining in dry air, after nitrogen, oxygen, argon and carbon dioxide were removed. Neon was the second of these three rare gases to be discovered and was immediately recognized as a new element from its bright red emission spectrum. The name neon is derived from the Greek word, νέον, neuter singular form of νέος (neos), meaning new. Neon is chemically inert, and no uncharged neon compounds are known. The compounds of neon currently known include ionic molecules, molecules held together by van der Waals forces and clathrates.

During cosmic nucleogenesis of the elements, large amounts of neon are built up from the alpha-capture fusion process in stars. Although neon is a very common element in the universe and solar system (it is fifth in cosmic abundance after hydrogen, helium, oxygen and carbon), it is rare on Earth. It composes about 18.2 ppm of air by volume (this is about the same as the molecular or mole fraction) and a smaller fraction in Earth's crust. The reason for neon's relative scarcity on Earth and the inner (terrestrial) planets is that neon is highly volatile and forms no compounds to fix it to solids. As a result, it escaped from the planetesimals under the warmth of the newly ignited Sun in the early Solar System. Even the outer atmosphere of Jupiter is somewhat depleted of neon, although for a different reason.Neon gives a distinct reddish-orange glow when used in low-voltage neon glow lamps, high-voltage discharge tubes and neon advertising signs. The red emission line from neon also causes the well known red light of helium–neon lasers. Neon is used in some plasma tube and refrigerant applications but has few other commercial uses. It is commercially extracted by the fractional distillation of liquid air. Since air is the only source, it is considerably more expensive than helium.

Neon-burning process

The neon-burning process (nuclear decay) is a set of nuclear fusion reactions that take place in massive stars (at least 8 Solar masses). Neon burning requires high temperatures and densities (around 1.2×109 K or 100 KeV and 4×109 kg/m3).

At such high temperatures photodisintegration becomes a significant effect, so some neon nuclei decompose, releasing alpha particles:

Alternatively:

where the neutron consumed in the first step is regenerated in the second.

Neon burning takes place after carbon burning has consumed all carbon in the core and built up a new oxygen-neon-sodium-magnesium core. The core ceases producing fusion energy and contracts. This contraction increases density and temperature up to the ignition point of neon burning. The increased temperature around the core allows carbon to burn in a shell, and there will be shells burning helium and hydrogen outside.

During neon burning, oxygen and magnesium accumulate in the central core while neon is consumed. After a few years the star consumes all its neon and the core ceases producing fusion energy and contracts. Again, gravitational pressure takes over and compresses the central core, increasing its density and temperature until the oxygen-burning process can start.

Outline of astronomy

The following outline is provided as an overview of and topical guide to astronomy:

Astronomy – studies the universe beyond Earth, including its formation and development, and the evolution, physics, chemistry, meteorology, and motion of celestial objects (such as galaxies, planets, etc.) and phenomena that originate outside the atmosphere of Earth (such as the cosmic background radiation).

Oxygen-burning process

The oxygen-burning process is a set of nuclear fusion reactions that take place in massive stars that have used up the lighter elements in their cores. Oxygen-burning is preceded by the neon-burning process and succeeded by the silicon-burning process. As the neon-burning process ends, the core of the star contracts and heats until it reaches the ignition temperature for oxygen burning. Oxygen burning reactions are similar to those of carbon burning; however, they must occur at higher temperatures and densities due to the larger Coulomb barrier of oxygen. Oxygen in the core ignites in the temperature range of (1.5–2.6)×109 K and in the density range of (2.6–6.7)×109g/cm3. The principal reactions are given below, where the branching ratios assume that the deuteron channel is open (at high temperatures):

Near 2×109K, the oxygen burning reaction rate is approximately 2.8×10−12(T9/2)33, where T9 is the temperature in billions of Kelvin. Overall, the major products of the oxygen-burning process are 28Si, 32,33,34S, 35,37Cl, 36,38Ar, 39,41K, and 40,42Ca. Of these, 28Si and 32S constitute 90% of the final composition. The oxygen fuel within the core of the star is exhausted after 0.01–5 years depending on the star's mass and other parameters. The silicon-burning process which follows creates iron, but this iron cannot react further to create energy to support the star.

During the oxygen-burning process, proceeding outward, there is an oxygen-burning shell, followed by a neon shell, a carbon shell, a helium shell, and a hydrogen shell. The oxygen-burning process is the last nuclear reaction in the star's core which does not proceed via the alpha process.

Sodium

Sodium is a chemical element with the symbol Na (from Latin natrium) and atomic number 11. It is a soft, silvery-white, highly reactive metal. Sodium is an alkali metal, being in group 1 of the periodic table, because it has a single electron in its outer shell, which it readily donates, creating a positively charged ion—the Na+ cation. Its only stable isotope is 23Na. The free metal does not occur in nature, and must be prepared from compounds. Sodium is the sixth most abundant element in the Earth's crust and exists in numerous minerals such as feldspars, sodalite, and rock salt (NaCl). Many salts of sodium are highly water-soluble: sodium ions have been leached by the action of water from the Earth's minerals over eons, and thus sodium and chlorine are the most common dissolved elements by weight in the oceans.

Sodium was first isolated by Humphry Davy in 1807 by the electrolysis of sodium hydroxide. Among many other useful sodium compounds, sodium hydroxide (lye) is used in soap manufacture, and sodium chloride (edible salt) is a de-icing agent and a nutrient for animals including humans.

Sodium is an essential element for all animals and some plants. Sodium ions are the major cation in the extracellular fluid (ECF) and as such are the major contributor to the ECF osmotic pressure and ECF compartment volume. Loss of water from the ECF compartment increases the sodium concentration, a condition called hypernatremia. Isotonic loss of water and sodium from the ECF compartment decreases the size of that compartment in a condition called ECF hypovolemia.

By means of the sodium-potassium pump, living human cells pump three sodium ions out of the cell in exchange for two potassium ions pumped in; comparing ion concentrations across the cell membrane, inside to outside, potassium measures about 40:1, and sodium, about 1:10. In nerve cells, the electrical charge across the cell membrane enables transmission of the nerve impulse—an action potential—when the charge is dissipated; sodium plays a key role in that activity.

Star

A star is an astronomical object consisting of a luminous spheroid of plasma held together by its own gravity. The nearest star to Earth is the Sun. Many other stars are visible to the naked eye from Earth during the night, appearing as a multitude of fixed luminous points in the sky due to their immense distance from Earth. Historically, the most prominent stars were grouped into constellations and asterisms, the brightest of which gained proper names. Astronomers have assembled star catalogues that identify the known stars and provide standardized stellar designations. However, most of the estimated 300 sextillion (3×1023) stars in the Universe are invisible to the naked eye from Earth, including all stars outside our galaxy, the Milky Way.

For at least a portion of its life, a star shines due to thermonuclear fusion of hydrogen into helium in its core, releasing energy that traverses the star's interior and then radiates into outer space. Almost all naturally occurring elements heavier than helium are created by stellar nucleosynthesis during the star's lifetime, and for some stars by supernova nucleosynthesis when it explodes. Near the end of its life, a star can also contain degenerate matter. Astronomers can determine the mass, age, metallicity (chemical composition), and many other properties of a star by observing its motion through space, its luminosity, and spectrum respectively. The total mass of a star is the main factor that determines its evolution and eventual fate. Other characteristics of a star, including diameter and temperature, change over its life, while the star's environment affects its rotation and movement. A plot of the temperature of many stars against their luminosities produces a plot known as a Hertzsprung–Russell diagram (H–R diagram). Plotting a particular star on that diagram allows the age and evolutionary state of that star to be determined.

A star's life begins with the gravitational collapse of a gaseous nebula of material composed primarily of hydrogen, along with helium and trace amounts of heavier elements. When the stellar core is sufficiently dense, hydrogen becomes steadily converted into helium through nuclear fusion, releasing energy in the process. The remainder of the star's interior carries energy away from the core through a combination of radiative and convective heat transfer processes. The star's internal pressure prevents it from collapsing further under its own gravity. A star with mass greater than 0.4 times the Sun's will expand to become a red giant when the hydrogen fuel in its core is exhausted. In some cases, it will fuse heavier elements at the core or in shells around the core. As the star expands it throws a part of its mass, enriched with those heavier elements, into the interstellar environment, to be recycled later as new stars. Meanwhile, the core becomes a stellar remnant: a white dwarf, a neutron star, or, if it is sufficiently massive, a black hole.

Binary and multi-star systems consist of two or more stars that are gravitationally bound and generally move around each other in stable orbits. When two such stars have a relatively close orbit, their gravitational interaction can have a significant impact on their evolution. Stars can form part of a much larger gravitationally bound structure, such as a star cluster or a galaxy.

Stellar nucleosynthesis

Stellar nucleosynthesis is the theory explaining the creation (nucleosynthesis) of chemical elements by nuclear fusion reactions between atoms within stars. Stellar nucleosynthesis has occurred continuously since the original creation of hydrogen, helium and lithium during the Big Bang. It is a highly predictive theory that today yields excellent agreement between calculations based upon it and the observed abundances of the elements. It explains why the observed abundances of elements in the universe grow over time and why some elements and their isotopes are much more abundant than others. The theory was initially proposed by Fred Hoyle in 1946, who later refined it in 1954. Further advances were made, especially to nucleosynthesis by neutron capture of the elements heavier than iron, by Margaret Burbidge, Geoffrey Burbidge, William Alfred Fowler and Hoyle in their famous 1957 B2FH paper, which became one of the most heavily cited papers in astrophysics history.

Stars evolve because of changes in their composition (the abundance of their constituent elements) over their lifespans, first by burning hydrogen (main sequence star), then helium (red giant star), and progressively burning higher elements. However, this does not by itself significantly alter the abundances of elements in the universe as the elements are contained within the star. Later in its life, a low-mass star will slowly eject its atmosphere via stellar wind, forming a planetary nebula, while a higher–mass star will eject mass via a sudden catastrophic event called a supernova. The term supernova nucleosynthesis is used to describe the creation of elements during the evolution and explosion of a pre-supernova massive star (12–35 times the mass of the sun). Those massive stars are the most prolific source of new isotopes from carbon (Z = 6) to nickel (Z = 28).

The advanced sequence of burning fuels is driven by gravitational collapse and its associated heating, resulting in the subsequent burning of carbon, oxygen and silicon. However, most of the nucleosynthesis in the mass range A = 28–56 (from silicon to nickel) is actually caused by the upper layers of the star collapsing onto the core, creating a compressional shock wave rebounding outward. The shock front briefly raises temperatures by roughly 50%, thereby causing furious burning for about a second. This final burning in massive stars, called explosive nucleosynthesis or supernova nucleosynthesis, is the final epoch of stellar nucleosynthesis.

A stimulus to the development of the theory of nucleosynthesis was the discovery of variations in the abundances of elements found in the universe. The need for a physical description was already inspired by the relative abundances of isotopes of the chemical elements in the solar system. Those abundances, when plotted on a graph as a function of atomic number of the element, have a jagged sawtooth shape that varies by factors of tens of millions (see history of nucleosynthesis theory). This suggested a natural process that is not random. A second stimulus to understanding the processes of stellar nucleosynthesis occurred during the 20th century, when it was realized that the energy released from nuclear fusion reactions accounted for the longevity of the Sun as a source of heat and light.

Type II supernova

A Type II supernova (plural: supernovae or supernovas) results from the rapid collapse and violent explosion of a massive star. A star must have at least 8 times, but no more than 40 to 50 times, the mass of the Sun (M☉) to undergo this type of explosion. Type II supernovae are distinguished from other types of supernovae by the presence of hydrogen in their spectra. They are usually observed in the spiral arms of galaxies and in H II regions, but not in elliptical galaxies.

Stars generate energy by the nuclear fusion of elements. Unlike the Sun, massive stars possess the mass needed to fuse elements that have an atomic mass greater than hydrogen and helium, albeit at increasingly higher temperatures and pressures, causing increasingly shorter stellar life spans. The degeneracy pressure of electrons and the energy generated by these fusion reactions are sufficient to counter the force of gravity and prevent the star from collapsing, maintaining stellar equilibrium. The star fuses increasingly higher mass elements, starting with hydrogen and then helium, progressing up through the periodic table until a core of iron and nickel is produced. Fusion of iron or nickel produces no net energy output, so no further fusion can take place, leaving the nickel–iron core inert. Due to the lack of energy output creating outward thermal pressure, the core contracts due to gravity until the overlying weight of the star can be supported largely by electron degeneracy pressure.

When the compacted mass of the inert core exceeds the Chandrasekhar limit of about 1.4 M☉, electron degeneracy is no longer sufficient to counter the gravitational compression. A cataclysmic implosion of the core takes place within seconds. Without the support of the now-imploded inner core, the outer core collapses inwards under gravity and reaches a velocity of up to 23% of the speed of light and the sudden compression increases the temperature of the inner core to up to 100 billion kelvins. Neutrons and neutrinos are formed via reversed beta-decay, releasing about 1046 joules (100 foe) in a ten-second burst. Also, the collapse of the inner core is halted by neutron degeneracy, causing the implosion to rebound and bounce outward. The energy of this expanding shock wave is sufficient to disrupt the overlying stellar material and accelerate it to escape velocity, forming a supernova explosion. The shock wave and extremely high temperature and pressure rapidly dissipate but are present for long enough to allow for a brief period during which the

production of elements heavier than iron occurs. Depending on initial size of the star, the remnants of the core form a neutron star or a black hole. Because of the underlying mechanism, the resulting supernova is also described as a core-collapse supernova.

There exist several categories of Type II supernova explosions, which are categorized based on the resulting light curve—a graph of luminosity versus time—following the explosion. Type II-L supernovae show a steady (linear) decline of the light curve following the explosion, whereas Type II-P display a period of slower decline (a plateau) in their light curve followed by a normal decay. Type Ib and Ic supernovae are a type of core-collapse supernova for a massive star that has shed its outer envelope of hydrogen and (for Type Ic) helium. As a result, they appear to be lacking in these elements.

Radioactive
decay
Stellar
nucleosynthesis
Other
processes

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