View source for Stellar nucleosynthesis

{{Short description|Creation of chemical elements within stars}}
[[File:Nuclear energy generation.svg|right|upright=1.5|thumb|[[Logarithmic scale]] plot of the relative energy output ({{mvar|ε}}) of the following fusion processes at different temperatures ({{mvar|T}}):
{{legend-line|solid lime 2px|[[Proton–proton chain]] (PP)}}
{{legend-line|solid blue 2px|[[CNO cycle]]}}
{{legend-line|solid red 2px|[[Triple-alpha process|Triple-α process]]}}
{{legend-line|dashed black 2px|Combined energy generation of PP and CNO within a star}}
{{legend-line|solid orange 1px|The Sun's core temperature, at which PP is more efficient}}]]

In [[astrophysics]], '''stellar nucleosynthesis''' is the [[nucleosynthesis|creation]] of [[chemical element]]s by [[nuclear fusion]] reactions within [[star]]s. Stellar nucleosynthesis has occurred since the [[Big Bang nucleosynthesis|original creation]] of [[hydrogen]], [[helium]] and [[lithium]] during the [[Big Bang]]. As a [[predictive power|predictive theory]], it yields accurate estimates of the [[Abundance of the chemical elements|observed abundances]] of the elements. It explains why the observed abundances of elements change over time and why some elements and their [[isotope]]s are much more abundant than others. The theory was initially proposed by [[Fred Hoyle]] in 1946,<ref name=Hoyle1946/> who later refined it in 1954.<ref name=Hoyle1954>
{{cite journal
 |last1=Hoyle |first1=F.
 |year=1954
 |title=On Nuclear Reactions Occurring in Very Hot STARS. I. The Synthesis of Elements from Carbon to Nickel
 |journal=[[The Astrophysical Journal Supplement Series]]
 |volume=1 |page=121
 |bibcode=1954ApJS....1..121H
 |doi=10.1086/190005
}}</ref> Further advances were made, especially to nucleosynthesis by [[neutron capture]] of the elements heavier than [[iron]], by [[Margaret Burbidge|Margaret]] and [[Geoffrey Burbidge]], [[William Alfred Fowler]] and [[Fred Hoyle]] in their famous 1957 [[B2FH paper|B<sup>2</sup>FH paper]],<ref name=B2FH>
{{cite journal
 |last1=Burbidge |first1=E. M.
 |last2=Burbidge |first2=G. R.
 |last3=Fowler |first3=W.A.
 |last4=Hoyle |first4=F.
 |year=1957
 |title=Synthesis of the Elements in Stars
 |url=https://authors.library.caltech.edu/45747/1/BURrmp57.pdf
 |journal=[[Reviews of Modern Physics]]
 |volume=29 |issue=4 |pages=547–650
 |bibcode=1957RvMP...29..547B
 |doi=10.1103/RevModPhys.29.547
|doi-access=free
 }}</ref> which became one of the most heavily cited papers in astrophysics history.

[[Stellar evolution|Stars evolve]] because of changes in their composition (the abundance of their constituent elements) over their lifespans, first by [[hydrogen burning|burning hydrogen]] ([[main sequence]] star), then [[helium burning|helium]] ([[horizontal branch]] star), and progressively burning [[triple-alpha process|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 explosion of a massive star or [[white dwarf]].

The advanced sequence of burning fuels is driven by [[gravitational collapse]] and its associated heating, resulting in the subsequent burning of [[carbon burning|carbon]], [[Oxygen-burning process|oxygen]] and [[silicon burning|silicon]]. However, most of the nucleosynthesis in the mass range {{nowrap|1=''[[mass number|A]]'' = 28–56}} (from silicon to nickel) is actually caused by the upper layers of the star [[Supernova#Core collapse|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 [[Abundance of the chemical elements#Universe|abundances of elements found in the universe]]. The need for a physical description was already inspired by the relative abundances of the chemical elements in the solar system. Those abundances, when plotted on a graph as a function of the atomic number of the element, have a jagged sawtooth shape that varies by factors of tens of millions (see [[history of nucleosynthesis theory]]).<ref>
{{cite journal
 |last1=Suess |first1=H. E.
 |last2=Urey |first2=H. C.
 |year=1956
 |title=Abundances of the Elements
 |journal=[[Reviews of Modern Physics]]
 |volume=28 |issue=1 |pages=53–74
 |bibcode=1956RvMP...28...53S
 |doi=10.1103/RevModPhys.28.53
}}</ref> 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.<ref name=Clayton1968>
{{cite book
 |last=Clayton |first=D. D.
 |year=1968
 |title=Principles of Stellar Evolution and Nucleosynthesis
 |publisher=University of Chicago Press
}}</ref>

==History==
[[File:Arthur Stanley Eddington.jpg|250px|right|thumb|In 1920, [[Arthur Eddington]] proposed that stars obtained their energy from [[Nuclear fusion#Astrophysical reaction chains|nuclear fusion]] of [[hydrogen]] to form [[helium]] and also raised the possibility that the heavier elements are produced in stars.]]
In 1920, [[Arthur Eddington]], on the basis of the precise measurements of atomic masses by [[Francis William Aston|F.W. Aston]] and a preliminary suggestion by [[Jean Perrin]], proposed that stars obtained their energy from [[nuclear fusion]] of [[hydrogen]] to form [[helium]] and raised the possibility that the heavier elements are produced in stars.<ref>
{{cite journal
 |last1=Eddington |first1=A. S.
 |year=1920
 |title=The internal constitution of the stars
 |journal=[[The Observatory (journal)|The Observatory]]
 |volume=43 |issue=1341
 |pages=341–358
 |doi=10.1126/science.52.1341.233
 |pmid=17747682
 |bibcode=1920Obs....43..341E
|url=https://zenodo.org/record/1429642
 }}</ref><ref>
{{cite journal
 |last1=Eddington |first1=A. S.
 |year=1920
 |title=The Internal Constitution of the Stars
 |journal=[[Nature (journal)|Nature]]
 |volume=106 |issue=2653 |pages=233–240
 |bibcode=1920Natur.106...14E
 |doi=10.1038/106014a0
|pmid=17747682
 |url=https://zenodo.org/record/1429642
 |doi-access=free
 }}</ref><ref>
{{cite web
 |last=Selle |first=D.
 |date=October 2012
 |title=Why the Stars Shine
 |url=http://www.astronomyhouston.org/sites/default/files/guidestar/2012October.pdf
 |work=Guidestar
 |pages=6–8
 |publisher=Houston Astronomical Society
 |archive-url=https://web.archive.org/web/20131203024638/http://www.astronomyhouston.org/sites/default/files/guidestar/2012October.pdf
 |archive-date=2013-12-03
 |url-status=live
}}</ref> This was a preliminary step toward the idea of stellar nucleosynthesis. In 1928 [[George Gamow]] derived what is now called the [[Gamow factor]], a [[quantum mechanics|quantum-mechanical]] formula yielding the probability for two contiguous nuclei to overcome the electrostatic [[Coulomb barrier]] between them and approach each other closely enough to undergo nuclear reaction due to the [[strong nuclear force]] which is effective only at very short distances.<ref>Krane, K. S., ''Modern Physics'' ([[Hoboken, New Jersey|Hoboken, NJ]]: [[Wiley (publisher)|Wiley]], 1983), [https://books.google.com/books?id=-x-VDwAAQBAJ&pg=PA410 p. 410].</ref>{{rp|410}} In the following decade the Gamow factor was used by [[Robert d'Escourt Atkinson|Atkinson]] and [[Fritz Houtermans|Houtermans]] and later by [[Edward Teller]] and Gamow himself to derive the rate at which nuclear reactions would occur at the high temperatures believed to exist in stellar interiors.

In 1939, in a [[Nobel Prize#Nobel lecture|Nobel lecture]] entitled "Energy Production in Stars", [[Hans Bethe]] analyzed the different possibilities for reactions by which hydrogen is fused into helium.<ref>
{{cite journal
 |last1=Bethe |first1=H. A.
 |year=1939
 |title=Energy Production in Stars
 |journal=[[Physical Review]]
 |volume=55 |issue=5 |pages=434–456
 |bibcode=1939PhRv...55..434B
 |doi=10.1103/PhysRev.55.434
|pmid=17835673
 |doi-access=free
 }}</ref> He defined two processes that he believed to be the sources of energy in stars. The first one, the [[proton–proton chain reaction]], is the dominant energy source in stars with masses up to about the mass of the Sun. The second process, the [[CNO cycle|carbon–nitrogen–oxygen cycle]], which was also considered by [[Carl Friedrich von Weizsäcker]] in 1938, is more important in more massive main-sequence stars.<ref>{{cite book|last=Lang|first=K. R.|title=The Life and Death of Stars|year=2013|publisher=Cambridge University Press|isbn=978-1-107-01638-5|page=[https://books.google.com/books?id=MN-UCkUK9pcC&pg=PA167 167]}}.</ref>{{rp|167}} These works concerned the energy generation capable of keeping stars hot. A clear physical description of the proton–proton chain and of the CNO cycle appears in a 1968 textbook.<ref>Clayton, D. D. (1968). ''Principles of Stellar Evolution and Nucleosynthesis''. [[University of Chicago Press]]. [https://books.google.com/books?id=8HSGFThnbvkC&pg=PT365 p. 365].</ref>{{rp|365}} Bethe's two papers did not address the creation of heavier nuclei, however. That theory was begun by Fred Hoyle in 1946 with his argument that a collection of very hot nuclei would assemble thermodynamically into [[iron]].<ref name=Hoyle1946>
{{cite journal
 |last=Hoyle |first=F.
 |year=1946
 |title=The synthesis of the elements from hydrogen
 |journal=[[Monthly Notices of the Royal Astronomical Society]]
 |volume=106 |issue=5 |pages=343–383
 |bibcode=1946MNRAS.106..343H
 |doi=10.1093/mnras/106.5.343
|doi-access=free
 }}</ref> Hoyle followed that in 1954 with a paper describing how advanced fusion stages within massive stars would synthesize the elements from carbon to iron in mass.<ref name=Hoyle1954/><ref>
{{cite journal
 |last1=Clayton |first1=D. D.
 |title=History of Science: Hoyle's Equation
 |journal=[[Science (journal)|Science]]
 |year=2007
 |volume=318 |issue=5858 |pages=1876–1877
 |doi=10.1126/science.1151167
|pmid=18096793
 |s2cid=118423007
 }}</ref>

Hoyle's theory was extended to other processes, beginning with the publication of the 1957 review paper "Synthesis of the Elements in Stars" by [[Margaret Burbidge|Burbidge]], [[Geoffrey Burbidge|Burbidge]], [[William Alfred Fowler|Fowler]] and [[Fred Hoyle|Hoyle]], more commonly referred to as the [[B2FH paper|B<sup>2</sup>FH paper]].<ref name=B2FH/> This review paper collected and refined earlier research into a heavily cited picture that gave promise of accounting for the observed relative abundances of the elements; but it did not itself enlarge Hoyle's 1954 picture for the origin of primary nuclei as much as many assumed, except in the understanding of nucleosynthesis of those elements heavier than iron by neutron capture. Significant improvements were made by [[Alastair G. W. Cameron]] and by [[Donald D. Clayton]]. In 1957 Cameron presented his own independent approach to nucleosynthesis,<ref>{{cite report
 |last=Cameron |first=A. G. W.
 |year=1957
 |title=Stellar Evolution, Nuclear Astrophysics, and Nucleogenesis
 |url=https://fas.org/sgp/eprint/CRL-41.pdf
 |publisher=[[Atomic Energy of Canada Limited]]
 |id=Report CRL-41
}}</ref> informed by Hoyle's example, and introduced computers into time-dependent calculations of evolution of nuclear systems. Clayton calculated the first time-dependent models of the [[S-process|''s''-process]] in 1961<ref>
{{cite journal
 |last1=Clayton |first1=D. D.
 |last2=Fowler |first2=W. A.
 |last3=Hull |first3=T. E.
 |last4=Zimmerman |first4=B. A.
 |year=1961
 |title=Neutron capture chains in heavy element synthesis
 |journal=[[Annals of Physics]]
 |volume=12 |issue=3 |pages=331–408
 |bibcode=1961AnPhy..12..331C
 |doi=10.1016/0003-4916(61)90067-7
}}</ref> and of the [[R-process|''r''-process]] in 1965,<ref name=Seeger1965>
{{cite journal
 |last1=Seeger |first1=P. A.
 |last2=Fowler |first2=W. A.
 |last3=Clayton |first3=D. D.
 |year=1965
 |title=Nucleosynthesis of Heavy Elements by Neutron Capture
 |journal=[[The Astrophysical Journal Supplement Series]]
 |volume=11 |pages=121–126
 |bibcode=1965ApJS...11..121S
 |doi=10.1086/190111
|url=https://tigerprints.clemson.edu/cgi/viewcontent.cgi?article=1307&context=physastro_pubs
 }}</ref> as well as of the burning of silicon into the abundant alpha-particle nuclei and iron-group elements in 1968,<ref name=Bodansky1968a>
{{cite journal
 |last1=Bodansky |first1=D.
 |last2=Clayton |first2=D. D.
 |last3=Fowler |first3=W. A.
 |year=1968
 |title=Nucleosynthesis During Silicon Burning
 |journal=[[Physical Review Letters]]
 |volume=20 |issue=4 |pages=161–164
 |bibcode=1968PhRvL..20..161B
 |doi=10.1103/PhysRevLett.20.161
|url=https://tigerprints.clemson.edu/cgi/viewcontent.cgi?article=1393&context=physastro_pubs
 }}</ref><ref name=Bodansky1968b>
{{cite journal
 |last1=Bodansky |first1=D.
 |last2=Clayton |first2=D. D.
 |last3=Fowler |first3=W. A.
 |year=1968
 |title=Nuclear Quasi-Equilibrium during Silicon Burning
 |journal=[[The Astrophysical Journal Supplement Series]]
 |volume=16 |page=299
 |bibcode=1968ApJS...16..299B
 |doi=10.1086/190176
|url=https://tigerprints.clemson.edu/physastro_pubs/312
 }}</ref> and discovered radiogenic chronologies<ref>
{{cite journal
 |last1=Clayton |first1=D. D.
 |year=1964
 |title=Cosmoradiogenic Chronologies of Nucleosynthesis
 |journal=[[The Astrophysical Journal]]
 |volume=139 |page=637
 |bibcode=1964ApJ...139..637C
 |doi=10.1086/147791
|url=https://tigerprints.clemson.edu/cgi/viewcontent.cgi?article=1305&context=physastro_pubs
 }}</ref> for determining the age of the elements.

[[File:Nucleosynthesis in a star.gif|thumb|Cross section of a [[supergiant]] showing nucleosynthesis and elements formed.]]

==Key reactions==
[[File:Nucleosynthesis periodic table.svg|thumb|right|500px|A version of the periodic table indicating the origins – including stellar nucleosynthesis – of the elements.]]
The most important reactions in stellar nucleosynthesis:
* [[Hydrogen]] fusion:
** [[Deuterium fusion]]
** The [[proton–proton chain]]
** The [[CNO cycle|carbon–nitrogen–oxygen cycle]]
* [[Helium]] fusion:
** The [[triple-alpha process]]
** The [[alpha process]]
* Fusion of heavier elements:
** [[Lithium burning]]: a process found most commonly in [[brown dwarf]]s
** [[Carbon-burning process]]
** [[Neon-burning process]]
** [[Oxygen-burning process]]
** [[Silicon-burning process]]
* Production of elements heavier than [[iron]]:
** [[Neutron]] capture:
*** The [[r-process]]
*** The [[s-process]]
** [[Proton]] capture:
*** The [[rp-process]]
*** The [[p-process]]
** [[Photodisintegration]]

===Hydrogen fusion===
{{Redirect|Hydrogen burning|the combustion of hydrogen gas|Hydrogen#Combustion}}
{{Main|Proton–proton chain reaction|CNO cycle|Deuterium fusion}}
{{multiple image
| align     = right
| direction = vertical
| width     = 300
| image1    = Fusion in the Sun.svg
| caption1  = '''Proton–proton chain reaction'''
| image2    = CNO Cycle.svg
| caption2  = '''CNO-I cycle'''<br />The helium nucleus is released at the top-left step.
}}
Hydrogen fusion (nuclear fusion of four protons to form a [[helium-4]] nucleus<ref name=jones2009/>) is the dominant process that generates energy in the cores of [[main sequence|main-sequence]] stars. It is also called "hydrogen burning", which should not be confused with the [[Chemical reaction|chemical]] [[hydrogen#Combustion|combustion of hydrogen]] in an [[oxidizing]] atmosphere. There are two predominant processes by which stellar hydrogen fusion occurs: [[proton–proton chain]] and the carbon–nitrogen–oxygen (CNO) cycle. Ninety percent of all stars, with the exception of [[white dwarfs]], are fusing hydrogen by these two processes.<ref>Seeds, M. A., ''Foundations of Astronomy'' ([[Belmont, California|Belmont, CA]]: [[Cengage|Wadsworth Publishing Company]], 1986), p. 245.</ref>{{rp|245}}

In the cores of lower-mass main-sequence stars such as the [[Sun]], the dominant energy production process is the [[proton–proton chain reaction]]. This creates a helium-4 nucleus through a sequence of reactions that begin with the fusion of two protons to form a [[deuterium]] nucleus (one proton plus one neutron) along with an ejected positron and neutrino.<ref name=bohm_vitense1992/> In each complete fusion cycle, the proton–proton chain reaction releases about 26.2&nbsp;MeV.<ref name=bohm_vitense1992/> The proton–proton chain reaction cycle is relatively insensitive to temperature; a 10% rise of temperature would increase energy production by this method by 46%, hence, this hydrogen fusion process can occur in up to a third of the star's radius and occupy half the star's mass. For stars above 35% of the Sun's mass,<ref name=aaa496_3_787/> the [[energy flux]] toward the surface is sufficiently low and energy transfer from the core region remains by [[radiative heat transfer]], rather than by [[Convection (heat transfer)|convective heat transfer]].<ref name=deloore_doom1992/> As a result, there is little mixing of fresh hydrogen into the core or fusion products outward.

In higher-mass stars, the dominant energy production process is the [[CNO cycle]], which is a [[catalytic cycle]] that uses nuclei of carbon, nitrogen and oxygen as intermediaries and in the end produces a helium nucleus as with the proton–proton chain.<ref name=bohm_vitense1992/> During a complete CNO cycle, 25.0&nbsp;MeV of energy is released. The difference in energy production of this cycle, compared to the proton–proton chain reaction, is accounted for by the energy lost through [[neutrino]] emission.<ref name=bohm_vitense1992/> The CNO cycle is very temperature sensitive, a 10% rise of temperature would produce a 350% rise in energy production. About 90% of the CNO cycle energy generation occurs within the inner 15% of the star's mass, hence it is strongly concentrated at the core.<ref name=jeffrey2010/> This results in such an intense outward energy flux that [[convective]] energy transfer becomes more important than does [[radiative transfer]]. As a result, the core region becomes a [[convection zone]], which stirs the hydrogen fusion region and keeps it well mixed with the surrounding proton-rich region.<ref name=karttunen_oja2007/> This core convection occurs in stars where the CNO cycle contributes more than 20% of the total energy. As the star ages and the core temperature increases, the region occupied by the convection zone slowly shrinks from 20% of the mass down to the inner 8% of the mass.<ref name=jeffrey2010/> The Sun produces on the order of 1% of its energy from the CNO cycle.<ref>{{Cite web|title=Neutrinos yield first experimental evidence of catalyzed fusion dominant in many stars|url=https://phys.org/news/2020-11-neutrinos-yield-experimental-evidence-catalyzed.html|access-date=2020-11-26|website=phys.org|language=en}}</ref>{{efn|Particle physicist Andrea Pocar points out, "Confirmation of CNO burning in our sun, where it operates at only one percent, reinforces our confidence that we understand how stars work."}}<ref>[[Gregory Robert Choppin|Choppin, G. R.]], [[Jan-Olov Liljenzin|Liljenzin, J.-O.]], [[Jan Rydberg|Rydberg, J.]], & Ekberg, C., ''Radiochemistry and Nuclear Chemistry'' (Cambridge, MA: [[Academic Press]], 2013), [https://books.google.com/books?id=CN88gBPtiucC&pg=PA357&redir_esc=y#v=onepage&q&f=false p. 357].</ref>{{rp|357}}<ref>{{Cite journal|last1=Agostini|first1=M.|last2=Altenmüller|first2=K.|last3=Appel|first3=S.|last4=Atroshchenko|first4=V.|last5=Bagdasarian|first5=Z.|last6=Basilico|first6=D.|last7=Bellini|first7=G.|last8=Benziger|first8=J.|last9=Biondi|first9=R.|last10=Bravo|first10=D.|last11=Caccianiga|first11=B.|date=25 November 2020|title=Experimental evidence of neutrinos produced in the CNO fusion cycle in the Sun|url=https://www.nature.com/articles/s41586-020-2934-0|journal=Nature|language=en|volume=587|issue=7835|pages=577–582|doi=10.1038/s41586-020-2934-0|pmid=33239797|issn=1476-4687|arxiv=2006.15115|bibcode=2020Natur.587..577B|s2cid=227174644}}</ref>{{efn|"This result therefore paves the way toward a direct measurement of the solar metallicity using CNO neutrinos. Our findings quantify the relative contribution of CNO fusion in the Sun to be of the order of 1 per cent."—M. Agostini, et al.}}

The type of hydrogen fusion process that dominates in a star is determined by the temperature dependency differences between the two reactions. The proton–proton chain reaction starts at temperatures about {{val|4|e=6|ul=K}},<ref name=reid_hawley2005/> making it the dominant fusion mechanism in smaller stars. A self-maintaining CNO chain requires a higher temperature of approximately {{val|1.6|e=7|u=K}}, but thereafter it increases more rapidly in efficiency as the temperature rises, than does the proton–proton reaction.<ref name=salaris_cassini2005/> Above approximately {{val|1.7|e=7|u=K}}, the CNO cycle becomes the dominant source of energy. This temperature is achieved in the cores of main-sequence stars with at least 1.3 times the mass of the [[Sun]].<ref name=apj701_1_837/> The Sun itself has a core temperature of about {{val|1.57|e=7|u=K}}.<ref>Wolf, E. L., ''Physics and Technology of Sustainable Energy'' ([[Oxford]], [[Oxford University Press]], 2018), [https://books.google.com/books?id=BP9eDwAAQBAJ&pg=PA5&redir_esc=y#v=onepage&q&f=false p. 5].</ref>{{rp|5}} As a main-sequence star ages, the core temperature will rise, resulting in a steadily increasing contribution from its CNO cycle.<ref name=jeffrey2010/>

===Helium fusion===
{{Main|Triple-alpha process|Alpha process}}
Main sequence stars accumulate helium in their cores as a result of hydrogen fusion, but the core does not become hot enough to initiate helium fusion.  Helium fusion first begins when a star leaves the [[red giant branch]] after accumulating sufficient helium in its core to ignite it. In stars around the mass of the Sun, this begins at the tip of the red giant branch with a [[helium flash]] from a [[Degenerate matter|degenerate]] helium core, and the star moves to the [[horizontal branch]] where it burns helium in its core. More massive stars ignite helium in their core without a flash and execute a [[blue loop]] before reaching the [[asymptotic giant branch]]. Such a star initially moves away from the AGB toward bluer colours, then loops back again to what is called the [[Hayashi track]]. An important consequence of blue loops is that they give rise to classical [[Cepheid variable]]s, of central importance in determining distances in the [[Milky Way]] and to nearby galaxies.<ref>Karttunen, H., Kröger, P., Oja, H., Poutanen, M., & Donner, K. J., eds., ''Fundamental Astronomy'' ([[Berlin]]/[[Heidelberg]]: [[Springer Science+Business Media|Springer]], 1987), [https://books.google.com/books?id=DjeVdb0sLEAC&pg=PA250&redir_esc=y#v=onepage&q&f=false p. 250].</ref>{{rp|250}} Despite the name, stars on a blue loop from the red giant branch are typically not blue in colour but are rather yellow giants, possibly Cepheid variables. They fuse helium until the core is largely [[carbon]] and [[oxygen]]. The most massive stars become supergiants when they leave the main sequence and quickly start helium fusion as they become [[red supergiant]]s. After the helium is exhausted in the core of a star, helium fusion will continue in a shell around the carbon–oxygen core.<ref name=jones2009/><ref name=deloore_doom1992/>

In all cases, helium is fused to carbon via the triple-alpha process, i.e., three helium nuclei are transformed into carbon via [[Beryllium-8|<sup>8</sup>Be]].<ref>Rehder, D., ''Chemistry in Space: From Interstellar Matter to the Origin of Life'' ([[Weinheim]]: [[Wiley-VCH]], 2010), [https://books.google.com/books?id=baI91e8lgm0C&pg=PT30&redir_esc=y#v=onepage&q&f=false p. 30].</ref>{{rp|30}} This can then form oxygen, neon, and heavier elements via the alpha process. In this way, the alpha process preferentially produces elements with even numbers of protons by the capture of helium nuclei. Elements with odd numbers of protons are formed by other fusion pathways.<ref>[[Michael Perryman|Perryman, M.]], ''The Exoplanet Handbook'' (Cambridge: Cambridge University Press, 2011), [https://books.google.com/books?id=ngtmDwAAQBAJ&pg=PA398&redir_esc=y#v=onepage&q&f=false p. 398].</ref>{{rp|398}}

==Reaction rate==
The reaction rate density between species ''A'' and ''B'', having number densities ''n''<sub>''A'',''B''</sub>, is given by:
<math display="block">r = n_A \, n_B \, k </math>
where ''k'' is the [[reaction rate constant]] of each single elementary binary reaction composing the [[nuclear fusion]] process:
<math display="block">k = \langle \sigma(v)\,v \rangle</math>
here, ''σ''(''v'') is the cross-section at relative velocity ''v'', and averaging is performed over all velocities.

Semi-classically, the cross section is proportional to <math>\pi\,\lambda^2</math>, where <math>\lambda = h/p</math> is the [[matter wave|de Broglie wavelength]]. Thus semi-classically the cross section is proportional to <math display="inline">\frac{m}{E}</math>.

However, since the reaction involves [[quantum tunneling]], there is an exponential damping at low energies that depends on [[Gamow factor]] ''E''<sub>G</sub>, giving an [[Arrhenius equation]]:
<math display="block">\sigma(E) = \frac{S(E)}{E} e^{-\sqrt{\frac{E_\text{G}}{E}}}</math>
where ''S''(''E'') depends on the details of the nuclear interaction, and has the dimension of an energy multiplied for a cross section.

One then integrates over all energies to get the total reaction rate, using the [[Maxwell–Boltzmann distribution#Distribution for the energy|Maxwell–Boltzmann distribution]] and the relation:
<math display="block">\frac{r}{V} = n_A n_B \int_0^{\infty}\frac{S(E)}{E} \, e^{-\sqrt{\frac{E_\text{G}}{E}}} 2\sqrt{\frac{E}{\pi(kT)^3}} e^{-\frac{E}{kT}} \,\sqrt{\frac{2E}{m_\text{R}}}dE</math>
where <math>m_\text{R} = \frac{m_1 m_2}{m_1 + m_2}</math> is the [[reduced mass]].

Since this integration has an exponential damping at high energies of the form <math>\sim e^{-\frac{E}{kT}}</math> and at low energies from the Gamow factor, the integral almost vanished everywhere except around the peak, called '''Gamow peak''',<ref>Iliadis, C., ''Nuclear Physics of Stars'' (Weinheim: Wiley-VCH, 2015), [https://books.google.com/books?id=kLZNCAAAQBAJ&pg=PA185 p. 185].</ref>{{rp|185}} at ''E''<sub>0</sub>, where:
<math display="block">\frac{\partial}{\partial E} \left( -\sqrt{\frac{E_\text{G}}{E}} - \frac{E}{kT}\right) \, = \, 0</math>

Thus:
<math display="block">E_0 = \left(\frac{1}{2}kT \sqrt{E_\text{G}}\right)^\frac{2}{3}</math>

The exponent can then be approximated around ''E''<sub>0</sub> as:
<math display="block">e^{-\frac{E}{kT} - \sqrt{\frac{E_\text{G}}{E}}} \approx e^{-\frac{3E_0}{kT}} \exp\left(-\frac{(E - E_0)^2}{\frac{4}{3} E_0 kT}\right)</math>

And the reaction rate is approximated as:<ref>{{Cite web|url=http://zuserver2.star.ucl.ac.uk/~idh/PHAS2112/Lectures/Current/Part7.pdf|title=University College London astrophysics course: lecture 7 – Stars|access-date=May 8, 2020|archive-url=https://web.archive.org/web/20170115214447/http://zuserver2.star.ucl.ac.uk/~idh/PHAS2112/Lectures/Current/Part7.pdf|archive-date=January 15, 2017|url-status=dead}}</ref>
<math display="block">\frac{r}{V} \approx n_A \, n_B  \, \frac{4\sqrt{2}}{\sqrt{3 m_\text{R}}}\, \sqrt{E_0} \frac{S(E_0)}{kT} e^{-\frac{3E_0}{kT}} </math>

Values of ''S''(''E''<sub>0</sub>) are typically {{nowrap|10<sup>−3</sup> – 10<sup>3</sup> [[keV]]·[[barn (unit)|b]]}}, but are damped by a huge factor when involving a [[beta decay]], due to the relation between the intermediate bound state (e.g. [[diproton]]) [[half-life]] and the beta decay half-life, as in the [[proton–proton chain reaction]]. Note that typical core temperatures in [[main-sequence star]]s give ''kT'' of the order of keV.<ref>Maoz, D., ''Astrophysics in a Nutshell'' ([[Princeton, New Jersey|Princeton]]: [[Princeton University Press]], 2007), [https://assets.press.princeton.edu/chapters/s3-10772.pdf ch. 3].</ref>{{rp|ch. 3}}

Thus, the limiting reaction in the [[CNO cycle]], proton capture by {{nuclide|nitrogen|14|link=yes}}, has ''S''(''E''<sub>0</sub>) ~ ''S''(0) = 3.5{{nbsp}}keV·b, while the limiting reaction in the [[proton–proton chain reaction]], the creation of [[deuterium]] from two protons, has a much lower ''S''(''E''<sub>0</sub>) ~ ''S''(0) = 4×10<sup>−22</sup>{{nbsp}}keV·b.<ref>{{Cite journal|last1=Adelberger|first1=Eric G.|author1-link=Eric G. Adelberger|last2=Austin|first2=Sam M.|last3=Bahcall|first3=John N.|last4=Balantekin|first4=A. B.|author4-link=A. Baha Balantekin|last5=Bogaert|first5=Gilles|last6=Brown|first6=Lowell S.|author6-link=Lowell S. Brown|last7=Buchmann|first7=Lothar|last8=Cecil|first8=F. Edward|last9=Champagne|first9=Arthur E.|last10=de Braeckeleer|first10=Ludwig|last11=Duba|first11=Charles A.|date=1998-10-01|title=Solar fusion cross sections|journal=Reviews of Modern Physics|language=en|volume=70|issue=4|pages=1265–1291|doi=10.1103/RevModPhys.70.1265|arxiv=astro-ph/9805121|bibcode=1998RvMP...70.1265A|s2cid=16061677|issn=0034-6861}}</ref><ref>{{Cite journal |arxiv = 1004.2318|doi = 10.1103/RevModPhys.83.195|bibcode = 2011RvMP...83..195A|title = Solar fusion cross sections. II. Theppchain and CNO cycles|year = 2011|last1 = Adelberger|first1 = E. G.|journal = Reviews of Modern Physics|volume = 83|issue = 1|pages = 195–245|s2cid = 119117147}}</ref> Incidentally, since the former reaction has a much higher Gamow factor, and due to the relative [[abundance of elements]] in typical stars, the two reaction rates are equal at a temperature value that is within the core temperature ranges of main-sequence stars.<ref>Goupil, M., Belkacem, K., Neiner, C., Lignières, F., & Green, J. J., eds., ''Studying Stellar Rotation and Convection: Theoretical Background and Seismic Diagnostics'' (Berlin/Heidelberg: Springer, 2013), [https://books.google.com/books?id=ovO5BQAAQBAJ&pg=PA211 p. 211].</ref>

==References==
===Notes===
{{notelist}}
===Citations===
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}}

==Further reading==
* {{cite journal|last=Bethe|first=H. A.|date=1939|title=Energy Production in Stars|journal=[[Physical Review]]|volume=55|issue=1|pages=541–547|doi=10.1103/PhysRev.55.103|pmid=17835673|bibcode = 1939PhRv...55..103B |doi-access=free}}
* {{cite journal|last=Bethe|first=H. A.|date=1939|title=Energy Production in Stars|journal=[[Physical Review]]|volume=55|issue=5|pages=434–456|doi= 10.1103/PhysRev.55.434|pmid=17835673|bibcode = 1939PhRv...55..434B |doi-access=free}}
* {{cite journal|last=Hoyle|first=F.|date=1954|title=On Nuclear Reactions occurring in very hot stars: Synthesis of elements from carbon to nickel|journal=[[Astrophysical Journal Supplement]]|volume=1|pages=121–146|doi=10.1086/190005|bibcode = 1954ApJS....1..121H }}
* {{cite book|last=Clayton|first=Donald D.|author-link=Donald D. Clayton|title=Principles of Stellar Evolution and Nucleosynthesis|url=https://archive.org/details/principlesofstel00clay|url-access=registration|publisher=[[McGraw-Hill]]|location=New York|date=1968}}
* {{cite arXiv|eprint=astro-ph/0405568|last1=Ray|first1=A.|title=Stars as thermonuclear reactors: Their fuels and ashes|date=2004}}
* {{cite journal | display-authors= 6 | author= G. Wallerstein |author-link1=George Wallerstein| author2= I. Iben, Jr. |author-link2=Icko Iben| author3= P. Parker | author4= A. M. Boesgaard |author-link4=Ann Merchant Boesgaard| author5= G. M. Hale | author6= A. E. Champagne | author7= C. A. Barnes | author8= F. Käppeler | author9= V.V. Smith | author10= R. D. Hoffman | author11= F. X. Timmes | author12= C. Sneden | author13= R. N. Boyd | author14= B. S. Meyer | author15= D. L. Lambert |author-link15=David L. Lambert| title= Synthesis of the elements in stars: forty years of progress | journal= [[Reviews of Modern Physics]] | date= 1997 | volume= 69 | issue= 4 | pages= 995–1084 | url= http://cococubed.asu.edu/papers/wallerstein97.pdf | access-date= 2006-08-04 | doi= 10.1103/RevModPhys.69.995 | bibcode= 1997RvMP...69..995W | url-status= dead | archive-url= https://web.archive.org/web/20090326090810/http://cococubed.asu.edu/papers/wallerstein97.pdf | archive-date= 2009-03-26 | hdl= 2152/61093 | hdl-access= free }}
* {{cite journal | last=Woosley | first=S. E. |author2=A. Heger |author3=T. A. Weaver  | author1-link=Stanford E. Woosley| title=The evolution and explosion of massive stars | journal=[[Reviews of Modern Physics]] | date=2002 | volume=74 | issue=4 | pages=1015–1071 | bibcode=2002RvMP...74.1015W | doi=10.1103/RevModPhys.74.1015 | s2cid=55932331 }}
* {{cite book|last=Clayton|first=Donald D.|title=Handbook of Isotopes in the Cosmos|url=https://books.google.com/books?id=fXcdHyLUVnEC|publisher=[[Cambridge University Press]]|location=Cambridge|date=2003|isbn=978-0-521-82381-4}}
*{{cite book|title=Nuclear Physics of Stars |last=Iliadis |first=Christian |year=2007 |isbn=978-3-527-40602-9 |publisher=Wiley-VCH |location=Weinheim |url=https://onlinelibrary.wiley.com/doi/book/10.1002/9783527618750 |doi=10.1002/9783527618750}}

==External links==
* [https://www.nobelprize.org/prizes/themes/how-the-sun-shines/ "How the Sun Shines"], by [[John N. Bahcall]] (Nobel prize site, accessed 6 January 2020)
* [https://web.archive.org/web/19990129020628/http://helios.gsfc.nasa.gov/nucleo.html Nucleosynthesis] in [[NASA]]'s Cosmicopia

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