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| {{Redirect|Carbon burning|combustion of carbon containing compounds|combustion}}
| | I am 25 years old and my name is Denese Atkinson. I life in Honefoss (Norway).<br><br>Feel free to visit my web site ... [http://hemorrhoidtreatmentfix.com/hemorrhoid-relief hemorrhoid relief] |
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| The '''carbon-burning process''' or '''carbon fusion''' is a set of [[nuclear fusion]] reactions that take place in massive [[star]]s (at least 8 [[Solar mass|<math>\begin{smallmatrix}M_\odot\end{smallmatrix}</math>]] at birth) that have used up the lighter elements in their cores. It requires high temperatures (> 5×10<sup>8</sup> [[kelvin|K]] or 50 [[keV]]) and [[density|densities]] (> 3×10<sup>9</sup> kg/m<sup>3</sub>).<ref name=Ryan>{{cite book| author=Ryan, Sean G.; Norton, Andrew J. | title=Stellar Evolution and Nucleosynthesis | year=2010 | page=135| isbn=978-0-521-13320-3|publisher=Cambridge University Press|url=http://books.google.com/?id=PE4yGiU-JyEC&q=carbon+burning#v=onepage&q=carbong%20burning&f=false}}</ref>
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| 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.<ref name="Clayton">{{cite book|last=Clayton|first=Donald| url=http://books.google.com/books?id=8HSGFThnbvkC|title=Principles of Stellar Evolution and Nucleosynthesis|year=1983|publisher=University of Chicago Press|isbn=978-0-226-10953-4}}</ref> To get the right figures for a particular mass, and a particular stage of evolution, it is necessary to use a numerical [[stellar structure|stellar model]] computed with computer algorithms.<ref name="Siess">{{cite journal
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| | author=Siess L.
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| | title = Evolution of massive AGB stars. I. Carbon burning phase
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| | journal = Astronomy and Astrophysics
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| | year= 2007
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| | volume = 476
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| | issue=2
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| | pages = 893–909
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| | doi = 10.1051/0004-6361:20053043
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| | bibcode = 2006A&A...448..717S}}</ref> Such models are continually being refined based on [[particle 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).<ref>{{cite journal
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| | title = Rubidium-Rich Asymptotic Giant Branch Stars
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| | journal = Science
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| | eprint = arXiv:astro-ph/0611319
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| | year = 2006
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| | month = dec,
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| | volume = 314
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| | issue = 5806
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| | pages = 1751–1754
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| | doi = 10.1126/science.1133706
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| | bibcode = 2006Sci...314.1751G
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| | author = Hernandez, G. et al
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| | pmid = 17095658
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| |arxiv = astro-ph/0611319 }}</ref>
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| ==Fusion reactions==
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| The principal reactions are:<ref name="Camiel1992">{{cite book|last=Camiel|first=W. H.|coauthors=de Loore, C. Doom|chapter=Structure and evolution of single and binary stars|editor=Camiel W. H. de Loore|title=Volume 179 of Astrophysics and space science library|year=1992|publisher=Springer|isbn=978-0-7923-1768-5|url=http://books.google.com/?id=LJgNIi0vkeYC&dq=%22carbon+burning%22&q=carbon+burning#v=snippet&q=carbon%20burning&f=false|pages=95–97}}</ref>
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| <!-- Autogenerated using Phykiformulae 0.10 by [[User:SkyLined]]
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| C-12 + C-12 -> Ne-20 + He (+ 4.617 MeV)
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| C-12 + C-12 -> Na-23 + H (+ 2.241 MeV)
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| C-12 + C-12 -> Mg-23 + n (− 2.599 MeV)
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| 99 Alternatively:
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| C-12 + C-12 -> Mg-24 + y (+ 13.933 MeV)
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| C-12 + C-12 -> O-16 + 2 He 2 (− 0.113 MeV)
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| -->:{| border="0"
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| |- style="height:3em;"
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| ||[[Carbon-12|{{nuclide|carbon|12}}]] ||+ ||[[Carbon-12|{{nuclide|carbon|12}}]] ||→ ||[[Neon-20|{{nuclide|neon|20}}]] ||+ ||[[Helium|{{nuclide|helium|4}}]] ||+ ||4.617 [[electron volt|MeV]]
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| |- style="height:3em;"
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| |[[Carbon-12|{{nuclide|carbon|12}}]] ||+ ||[[Carbon-12|{{nuclide|carbon|12}}]] ||→ ||[[Sodium-23|{{nuclide|sodium|23}}]] ||+ ||[[Hydrogen|{{nuclide|hydrogen|1}}]] ||+ ||2.241 [[electron volt|MeV]]
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| |- style="height:3em;"
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| |[[Carbon-12|{{nuclide|carbon|12}}]] ||+ ||[[Carbon-12|{{nuclide|carbon|12}}]] ||→ ||[[Magnesium-23|{{nuclide|magnesium|23}}]] ||+ ||[[Neutron|n]] ||− ||2.599 [[electron volt|MeV]]
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| |- style="height:3em;"
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| |colspan=99|Alternatively:
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| |- style="height:3em;"
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| |[[Carbon-12|{{nuclide|carbon|12}}]] ||+ ||[[Carbon-12|{{nuclide|carbon|12}}]] ||→ ||[[Magnesium-24|{{nuclide|magnesium|24}}]] ||+ ||{{SubatomicParticle|link=yes|Gamma}} ||+ ||13.933 [[electron volt|MeV]]
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| |- style="height:3em;"
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| |[[Carbon-12|{{nuclide|carbon|12}}]] ||+ ||[[Carbon-12|{{nuclide|carbon|12}}]] ||→ ||[[Oxygen-16|{{nuclide|oxygen|16}}]] ||+ ||2 [[Helium|{{nuclide|helium|4}}]] ||colspan=2|− 0.113 [[electron volt|MeV]]
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| |}
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| ==Reaction products==
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| 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.<ref name=Rose1998>{{cite book|url=http://books.google.com/books?id=yaX0etDmbXMC|last=Rose,|first=William K.|title=Advanced Stellar Astrophysics|publisher=Cambridge University Press|year=1998|isbn=978-0-521-58833-1|pages=227–229}}</ref> 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.<ref name=Camiel1992/> 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]].<ref name=Rose1998-229>Rose (1998), pp. 229–234</ref>
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| 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 the electromagnetic interaction,<ref name=Camiel1992/> 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.
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| The last reaction is also very unlikely since it involves three reaction products,<ref name=Camiel1992/> 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.
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| 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.<ref name=Camiel1992/> In fact, a significant fraction of the Na-23 produced by the second reaction gets used up this way.<ref name=Rose1998/> 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, in stars between 9 and 11 [[solar mass]]es.<ref name=Ryan/><ref name=Camiel1992-97>Camiel (1992), pp.97–98</ref> So the end result of carbon burning is a mixture mainly of oxygen, neon, sodium and magnesium.<ref name="Siess"/><ref name=Camiel1992/>
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| 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.
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| The experimental and theoretical investigation of such resonances is still a subject of research.<ref>{{cite journal
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| | author = Strandberg, E. et al
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| | title = Mg24(α,γ)Si28 resonance parameters at low α-particle energies
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| | journal = Physical Review C
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| | date = May 2008
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| | volume = 77
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| | issue = 5
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| | pages = 055801-+
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| | doi = 10.1103/PhysRevC.77.055801
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| | bibcode = 2008PhRvC..77e5801S
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| }}</ref> A similar resonance increases the probability of the [[triple-alpha process]], which is responsible for the original production of carbon.
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| ==Neutrino losses==
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| [[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 [[electron]]s ([[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.
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| 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 10<sup>19</sup> pair productions<ref name="Clayton"/> 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.
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| 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.<ref name="Clayton"/> 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.
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| Up to helium burning, the neutrino losses are negligible, but from carbon burning the reduction in lifetime due to them roughly matches that 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 10<sup>7</sup> years, helium for 10<sup>6</sup> years and carbon for only 10<sup>3</sup> years.<ref name="WoosleyJanka">{{cite journal
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| | last=Woosley | first=S. | coauthors=Janka, H.-T.
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| | date=2006-01-12
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| | bibcode=2005NatPh...1..147W
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| | title=The Physics of Core-Collapse Supernovae
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| | journal=Nature Physics
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| | year=2005 | month=December
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| | volume=1 | issue=3 | pages=147–154
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| | doi=10.1038/nphys172|arxiv = astro-ph/0601261 }}</ref>
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| ==Stellar evolution==
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| {{main|Stellar evolution}}
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| 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.<ref name="Ostlie">Ostlie, Dale A. and Carrol, Bradley W., [http://www.amazon.com/dp/0805304029 ''An introduction to Modern Stellar Astrophysics''], Addison-Wesley (2007)</ref> 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.<ref name="Lecture19">Anderson, Scott R.,
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| [http://www.opencourse.info/astronomy/introduction/19.stars_death_high-mass/ ''Open Course: Astronomy: Lecture 19: Death of High-Mass Stars''], GEM (2001)</ref>
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| Stars with below 8–9 [[Solar mass]]es never reach high enough core temperature to burn carbon, instead ending their lives as carbon-oxygen [[white dwarf]]s after shell [[helium flash]]es gently expel the outer envelope in a [[planetary nebula]].<ref name="Siess"/><ref name=Ryan-147/>
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| In the late stages of carbon burning, stars with masses between 8 and 11 solar masses 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.<ref name="Siess"/> The core never reaches high enough temperature for further fusion burning of heavier elements than carbon.<ref name=Ryan-147/>
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| Stars with more than 11 solar masses proceed with the [[neon-burning process]] after contraction of the inert (O, Ne, Na, Mg) core raises the temperature sufficiently.<ref name=Ryan-147>Ryan (2010), pp.147–148</ref>
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| ==See also==
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| *[[Proton–proton chain reaction]]
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| *[[CNO process]]
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| *[[Triple alpha process]]
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| *[[Alpha process]]
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| *[[Carbon detonation]]
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| *[[Neon burning]]
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| ==References==
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| {{reflist|2}}
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| {{Nuclear processes}}
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| [[Category:Nucleosynthesis]]
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