Editing Stellar evolution (section)

==Mature stars==
[[File:Star types.svg|upright=1.35|thumb|Internal structures of [[main sequence|main-sequence stars]], convection zones with arrowed cycles and radiative zones with red flashes. To the left a '''low-mass''' [[red dwarf]], in the center a '''mid-sized''' [[Yellow dwarf star|yellow dwarf]] and at the right a '''massive''' blue-white main-sequence star.]]
Eventually the star's core exhausts its supply of hydrogen and the star begins to evolve off the [[main sequence]]. Without the outward [[radiation pressure]] generated by the fusion of hydrogen to counteract the force of [[gravity]], the core contracts until either [[electron degeneracy pressure]] becomes sufficient to oppose gravity or the core becomes hot enough (around 100 MK) for [[helium fusion]] to begin. Which of these happens first depends upon the star's mass.

===Low-mass stars===
What happens after a low-mass star ceases to produce energy through fusion has not been directly observed; the [[universe]] is around 13.8 billion years old, which is less time (by several orders of magnitude, in some cases) than it takes for fusion to cease in such stars.

Recent astrophysical models suggest that [[red dwarf]]s of {{Solar mass|0.1}} may stay on the main sequence for some six to twelve trillion years, gradually increasing in both [[temperature]] and [[luminosity]], and take several hundred billion years more to collapse, slowly, into a [[white dwarf]].<ref name="S&T 22">{{cite journal| title=Why the Smallest Stars Stay Small| journal=Sky & Telescope|date=November 1997| issue=22}}</ref><ref>{{cite journal| journal=Astronomische Nachrichten| volume= 326| issue=10| pages= 913–919| date= 2005| title=M dwarfs: planet formation and long term evolution| first=F. C.|last= Adams| author2= P. Bodenheimer| author3=G. Laughlin|bibcode=2005AN....326..913A|doi=10.1002/asna.200510440| doi-access=free}}</ref>  Such stars will not become red giants as the whole star is a [[convection zone]] and it will not develop a degenerate helium core with a shell burning hydrogen.  Instead, hydrogen fusion will proceed until almost the whole star is helium.

Slightly more [[massive star]]s do expand into [[red giant]]s, but their helium cores are not massive enough to reach the temperatures required for helium fusion so they never reach the tip of the red-giant branch.  When hydrogen shell burning finishes, these stars move directly off the red-giant branch like a post-[[Asymptotic giant branch|asymptotic-giant-branch]] (AGB) star, but at lower luminosity, to become a white dwarf.<ref name="endms" />  A star with an initial mass about {{Solar mass|0.6}} will be able to reach temperatures high enough to fuse helium, and these "mid-sized" stars go on to further stages of evolution beyond the red-giant branch.<ref name="lejeune">{{cite journal|last1=Lejeune|first1=T|last2=Schaerer|first2=D|year=2001|title=Database of Geneva stellar evolution tracks and isochrones for <math>(UBV)_\mathsf{J}(RI)_\mathsf{C} JHKLL'M</math>, HST-WFPC2, Geneva and Washington photometric systems|journal=Astronomy & Astrophysics|volume=366|issue=2|pages=538–546|bibcode=2001A&A...366..538L|doi=10.1051/0004-6361:20000214|arxiv=astro-ph/0011497|s2cid=6708419}}</ref>

===Mid-sized stars===
[[File:Evolutionary track 1m.svg|thumb|left|upright=1.35|The evolutionary track of a solar mass, solar metallicity, star from main sequence to post-AGB]]
Stars of roughly {{Solar mass|0.6–10}} become [[red giant]]s, which are large non-[[main sequence|main-sequence]] stars of [[stellar classification]] K or M. Red giants lie along the right edge of the Hertzsprung–Russell diagram due to their red color and large luminosity. Examples include [[Aldebaran]] in the constellation [[Taurus (constellation)|Taurus]] and [[Arcturus]] in the constellation of [[Boötes]].

Mid-sized stars are red giants during two different phases of their post-main-sequence evolution: red-giant-branch stars, with inert cores made of helium and hydrogen-burning shells, and asymptotic-giant-branch stars, with inert cores made of carbon and helium-burning shells inside the hydrogen-burning shells.<ref>{{harvtxt|Hansen|Kawaler|Trimble|2004|pages=55–56}}</ref>  Between these two phases, stars spend a period on the [[horizontal branch]] with a helium-fusing core.  Many of these helium-fusing stars cluster towards the cool end of the horizontal branch as K-type giants and are referred to as [[red clump]] giants.

====Subgiant phase====
{{Main|Subgiant}}
When a star exhausts the hydrogen in its core, it leaves the main sequence and begins to fuse hydrogen in a shell outside the core.  The core increases in mass as the shell produces more helium.  Depending on the mass of the helium core, this continues for several million to one or two billion years, with the star expanding and cooling at a similar or slightly lower luminosity to its main sequence state.  Eventually either the core becomes degenerate, in stars around the mass of the sun, or the outer layers cool sufficiently to become opaque, in more massive stars.  Either of these changes cause the hydrogen shell to increase in temperature and the [[luminosity]] of the star to increase, at which point the star expands onto the red-giant branch.<ref name="RyanNorton115">{{harvtxt|Ryan|Norton|2010|page=115}}</ref>

====Red-giant-branch phase====
[[Image:The life of Sun-like stars.jpg|thumb|upright=1.35|Artist's depiction of the life cycle of a Sun-like star, starting as a main-sequence star at lower left then expanding through the [[subgiant]] and [[giant star|giant]] phases, until its outer envelope is expelled to form a [[planetary nebula]] at upper right]]

{{Main|Red-giant branch}}
The expanding outer layers of the star are [[convection|convective]], with the material being mixed by turbulence from near the fusing regions up to the surface of the star.  For all but the lowest-mass stars, the fused material has remained deep in the stellar interior prior to this point, so the convecting envelope makes fusion products visible at the star's surface for the first time. At this stage of evolution, the results are subtle, with the largest effects, alterations to the [[isotopes]] of hydrogen and helium, being unobservable. The effects of the [[CNO cycle]] appear at the surface during the first [[dredge-up]], with lower <sup>12</sup>C/<sup>13</sup>C ratios and altered proportions of carbon and nitrogen. These are detectable with [[spectroscopy]] and have been measured for many evolved stars.

The helium core continues to grow on the red-giant branch.  It is no longer in thermal equilibrium, either degenerate or above the [[Schönberg–Chandrasekhar limit]], so it increases in temperature which causes the rate of fusion in the hydrogen shell to increase.  The star increases in luminosity towards the [[tip of the red-giant branch]].  Red-giant-branch stars with a degenerate helium core all reach the tip with very similar core masses and very similar luminosities, although the more massive of the red giants become hot enough to ignite helium fusion before that point.
====Horizontal branch====
{{Main|Horizontal branch|Red clump}}
In the helium cores of stars in the 0.6 to 2.0 solar mass range, which are largely supported by [[electron degeneracy pressure]], helium fusion will ignite on a timescale of days in a [[helium flash]]. In the nondegenerate cores of more massive stars, the ignition of helium fusion occurs relatively slowly with no flash.<ref>{{harvtxt|Ryan|Norton|2010|page=125}}</ref> The nuclear power released during the helium flash is very large, on the order of 10<sup>8</sup> times the [[Solar luminosity|luminosity of the Sun]] for a few days<ref name="RyanNorton115" /> and 10<sup>11</sup> times the luminosity of the Sun (roughly the luminosity of the [[Milky Way Galaxy]]) for a few seconds.<ref name="Prialnik151">{{harvtxt|Prialnik|2000|page=151}}</ref> However, the energy is consumed by the thermal expansion of the initially degenerate core and thus cannot be seen from outside the star.<ref name="RyanNorton115" /><ref name="Prialnik151" /><ref name="Deupree1996">{{cite journal|last1= Deupree|first1=R. G.|title= A Reexamination of the Core Helium Flash|journal= The Astrophysical Journal|volume=471|issue= 1|date= 1996-11-01|pages= 377–384|doi= 10.1086/177976|bibcode= 1996ApJ...471..377D|citeseerx= 10.1.1.31.44|s2cid=15585754 }}</ref> Due to the expansion of the core, the hydrogen fusion in the overlying layers slows and total energy generation decreases. The star contracts, although not all the way to the main sequence, and it migrates to the [[horizontal branch]] on the Hertzsprung–Russell diagram, gradually shrinking in radius and increasing its surface temperature.
[[File:The life cycle of a Sun-like star (annotated).jpg|thumb|upright=1.35|left|The change in size with time of a Sun-like star]]

Core helium flash stars evolve to the red end of the horizontal branch but do not migrate to higher temperatures before they gain a degenerate carbon-oxygen core and start helium shell burning.  These stars are often observed as a [[red clump]] of stars in the colour-magnitude diagram of a cluster, hotter and less luminous than the red giants. Higher-mass stars with larger helium cores move along the horizontal branch to higher temperatures, some becoming unstable pulsating stars in the yellow [[instability strip]] ([[RR Lyrae variables]]), whereas some become even hotter and can form a blue tail or blue hook to the horizontal branch. The morphology of the horizontal branch depends on parameters such as metallicity, age, and helium content, but the exact details are still being modelled.<ref name=parameters>{{Cite journal | last1 = Gratton | first1 = R. G. | last2 = Carretta | first2 = E. | last3 = Bragaglia | first3 = A. | last4 = Lucatello | first4 = S. | last5 = d'Orazi | first5 = V. | title = The second and third parameters of the horizontal branch in globular clusters | doi = 10.1051/0004-6361/200912572 | journal = Astronomy and Astrophysics | volume = 517 | pages = A81 | year = 2010 |arxiv = 1004.3862 |bibcode = 2010A&A...517A..81G | s2cid = 55701280 }}</ref>

====Asymptotic-giant-branch phase====
{{Main|Asymptotic giant branch}}
After a star has consumed the helium at the core, hydrogen and helium fusion continues in shells around a hot core of [[carbon]] and [[oxygen]]. The star follows the [[asymptotic giant branch]] on the Hertzsprung–Russell diagram, paralleling the original red-giant evolution, but with even faster energy generation (which lasts for a shorter time).<ref>{{Cite journal | last1 = Sackmann | first1 = I. -J. | last2 = Boothroyd | first2 = A. I. | last3 = Kraemer | first3 = K. E. | title = Our Sun. III. Present and Future | doi = 10.1086/173407 | journal = The Astrophysical Journal | volume = 418 | pages = 457 | year = 1993 |bibcode = 1993ApJ...418..457S | doi-access = free }}</ref>  Although helium is being burnt in a shell, the majority of the energy is produced by hydrogen burning in a shell further from the core of the star.  Helium from these hydrogen burning shells drops towards the center of the star and periodically the energy output from the helium shell increases dramatically.  This is known as a [[thermal pulse]] and they occur towards the end of the asymptotic-giant-branch phase, sometimes even into the post-asymptotic-giant-branch phase. Depending on mass and composition, there may be several to hundreds of thermal pulses.

There is a phase on the ascent of the asymptotic-giant-branch where a deep convective zone forms and can bring carbon from the core to the surface.  This is known as the second dredge up, and in some stars there may even be a third dredge up.  In this way a [[carbon star]] is formed, very cool and strongly reddened stars showing strong carbon lines in their spectra.  A process known as hot bottom burning may convert carbon into oxygen and nitrogen before it can be dredged to the surface, and the interaction between these processes determines the observed luminosities and spectra of carbon stars in particular clusters.<ref name=hbb>{{cite journal|author1=van Loon |author2=Zijlstra |author3=Whitelock|author4=Peter te Lintel Hekkert|author5=Chapman|author6=Cecile Loup|author7=Groenewegen|author8=Waters|author9=Trams|title=Obscured Asymptotic Giant Branch stars in the Magellanic Clouds IV. Carbon stars and OH/IR stars|date=1998|volume=329|issue=1|pages=169–85|journal= Monthly Notices of the Royal Astronomical Society|arxiv=astro-ph/9709119v1 |citeseerx=10.1.1.389.3269 |bibcode = 1996MNRAS.279...32Z |doi=10.1093/mnras/279.1.32 |doi-access=free |url=https://pure.uva.nl/ws/files/978044/1833_20357y.pdf}}</ref>

Another well known class of asymptotic-giant-branch stars is the [[Mira variable]]s, which pulsate with well-defined periods of tens to hundreds of days and large amplitudes up to about 10 magnitudes (in the visual, total luminosity changes by a much smaller amount). In more-massive stars the stars become more luminous and the pulsation period is longer, leading to enhanced mass loss, and the stars become heavily obscured at visual wavelengths.  These stars can be observed as [[OH/IR star]]s, pulsating in the infrared and showing OH [[maser]] activity.  These stars are clearly oxygen rich, in contrast to the carbon stars, but both must be produced by dredge ups.

====Post-AGB====
{{Main|Post-AGB}}
[[Image:NGC6543.jpg|thumb|The [[Cat's Eye Nebula]], a [[planetary nebula]] formed by the death of a star with about the same mass as the Sun]]
These mid-range stars ultimately reach the tip of the asymptotic-giant-branch and run out of fuel for shell burning. They are not sufficiently massive to start full-scale carbon fusion, so they contract again, going through a period of post-asymptotic-giant-branch superwind to produce a planetary nebula with an extremely hot central star. The central star then cools to a white dwarf. The expelled gas is relatively rich in heavy elements created within the star and may be particularly [[oxygen]] or [[carbon]] enriched, depending on the type of the star. The gas builds up in an expanding shell called a [[circumstellar envelope]] and cools as it moves away from the star, allowing [[Circumstellar dust|dust particles]] and molecules to form. With the high infrared energy input from the central star, ideal conditions are formed in these circumstellar envelopes for [[Astrophysical maser|maser]] excitation.

It is possible for thermal pulses to be produced once post-asymptotic-giant-branch evolution has begun, producing a variety of unusual and poorly understood stars known as born-again asymptotic-giant-branch stars.<ref name=bornagain>{{cite journal|bibcode=1991IAUS..145..363H|title=Atmospheres and Abundances of Blue Horizontal Branch Stars and Related Objects|journal=Evolution of Stars: The Photospheric Abundance Connection: Proceedings of the 145th Symposium of the International Astronomical Union|volume=145|pages=363|last1=Heber|first1=U.|year=1991}}</ref> These may result in extreme [[horizontal-branch]] stars ([[subdwarf B star]]s), hydrogen deficient post-asymptotic-giant-branch stars, variable planetary nebula central stars, and [[R Coronae Borealis variable]]s.

===Massive stars===
{{Main|Supergiant}}
[[Image:VLTI reconstructed view of the surface of Antares.jpg|thumb|left|Reconstructed image of [[Antares]], a red supergiant]]
In massive stars, the core is already large enough at the onset of the hydrogen burning shell that helium ignition will occur before electron degeneracy pressure has a chance to become prevalent. Thus, when these stars expand and cool, they do not brighten as dramatically as lower-mass stars; however, they were more luminous on the main sequence and they evolve to highly luminous supergiants.  Their cores become massive enough that they cannot support themselves by [[electron degeneracy]] and will eventually collapse to produce a [[neutron star]] or [[black hole]].{{Citation needed|date=May 2021}}

====Supergiant evolution====
Extremely massive stars (more than approximately {{Solar mass|40}}), which are very luminous and thus have very rapid stellar winds, lose mass so rapidly due to radiation pressure that they tend to strip off their own envelopes before they can expand to become [[red supergiant]]s, and thus retain extremely high surface temperatures (and blue-white color) from their main-sequence time onwards. The largest stars of the current generation are about {{Solar mass|100-150}} because the outer layers would be expelled by the extreme radiation. Although lower-mass stars normally do not burn off their outer layers so rapidly, they can likewise avoid becoming red giants or red supergiants if they are in binary systems close enough so that the companion star strips off the envelope as it expands, or if they rotate rapidly enough so that convection extends all the way from the core to the surface, resulting in the absence of a separate core and envelope due to thorough mixing.<ref>{{cite journal |first1=D. |last1=Vanbeveren |title=Massive stars |journal=The Astronomy and Astrophysics Review |date=1998 |volume=9 |issue=1–2 |pages=63–152 |doi=10.1007/s001590050015 |last2=De Loore |first2=C. |last3=Van Rensbergen |first3=W. |bibcode=1998A&ARv...9...63V|s2cid=189933559 }}</ref>

[[Image:Evolved star fusion shells.svg|right|thumb|The onion-like layers of a massive, evolved star just before core collapse (not to scale)]]
The core of a massive star, defined as the region depleted of hydrogen, grows hotter and denser as it accretes material from the fusion of hydrogen outside the core.  In sufficiently massive stars, the core reaches temperatures and densities high enough to fuse carbon and heavier elements via the [[alpha process]].  At the end of helium fusion, the core of a star consists primarily of carbon and oxygen.  In stars heavier than about {{solar mass|8}}, the carbon ignites and [[Carbon-burning process|fuses]] to form neon, sodium, and magnesium.  Stars somewhat less massive may partially ignite carbon, but they are unable to fully fuse the carbon before [[electron degeneracy]] sets in, and these stars will eventually leave an oxygen-neon-magnesium [[white dwarf]].<ref name=jones>{{cite journal |doi=10.1088/0004-637X/772/2/150 |title=Advanced Burning Stages and Fate of 8-10M☉Stars |journal=The Astrophysical Journal |volume=772 |issue=2 |pages=150 |year=2013 |last1=Jones |first1=S. |last2=Hirschi |first2=R. |last3=Nomoto |first3=K. |last4=Fischer |first4=T. |last5=Timmes |first5=F. X. |last6=Herwig |first6=F. |last7=Paxton |first7=B. |last8=Toki |first8=H. |last9=Suzuki |first9=T. |last10=Martínez-Pinedo |first10=G. |last11=Lam |first11=Y. H. |last12=Bertolli |first12=M. G. |arxiv=1306.2030 |bibcode=2013ApJ...772..150J |s2cid=118687195 }}</ref><ref name=woosley>{{cite journal |doi=10.1103/RevModPhys.74.1015 |title=The evolution and explosion of massive stars |journal=Reviews of Modern Physics |volume=74 |issue=4 |pages=1015–1071 |year=2002 |last1=Woosley |first1=S. E. |last2=Heger |first2=A. |last3=Weaver |first3=T. A. |bibcode=2002RvMP...74.1015W |s2cid=55932331 }}</ref>

The exact mass limit for full carbon burning depends on several factors such as metallicity and the detailed mass lost on the [[asymptotic giant branch]], but is approximately {{solar mass|8-9}}.<ref name=jones/>  After carbon burning is complete, the core of these stars reaches about {{solar mass|2.5}} and becomes hot enough for heavier elements to fuse.  Before oxygen starts to [[Oxygen-burning process|fuse]], neon begins to [[Electron capture|capture electrons]] which triggers [[Neon-burning process|neon burning]].  For a range of stars of approximately {{solar mass|8-12}}, this process is unstable and creates runaway fusion resulting in an [[electron capture supernova]].<ref name=nomoto1987>{{cite journal |author=Ken'ichi Nomoto |title=Evolution of 8–10 {{Solar mass}} stars toward electron capture supernovae. II – Collapse of an O + Ne + Mg core |journal=Astrophysical Journal |date=1987 |volume=322 |pages=206–214 |bibcode=1987ApJ...322..206N |doi=10.1086/165716}}</ref><ref name=woosley/>

In more massive stars, the fusion of neon proceeds without a runaway deflagration.  This is followed in turn by complete oxygen burning and [[Silicon-burning process|silicon burning]], producing a core consisting largely of [[iron-peak element]]s.  Surrounding the core are shells of lighter elements still undergoing fusion.  The timescale for complete fusion of a carbon core to an iron core is so short, just a few hundred years, that the outer layers of the star are unable to react and the appearance of the star is largely unchanged.  The iron core grows until it reaches an ''effective Chandrasekhar mass'', higher than the formal [[Chandrasekhar mass]] due to various corrections for the relativistic effects, entropy, charge, and the surrounding envelope.  The effective Chandrasekhar mass for an iron core varies from about {{solar mass|1.34}} in the least massive red supergiants to more than {{solar mass|1.8}} in more massive stars.  Once this mass is reached, electrons begin to be captured into the iron-peak nuclei and the core becomes unable to support itself.  The core collapses and the star is destroyed, either in a [[supernova]] or direct collapse to a [[black hole]].<ref name=woosley/>

====Supernova====
{{Main|Supernova}}
[[Image:Crab Nebula.jpg|thumb|left|The [[Crab Nebula]], the shattered remnants of a star which exploded as a supernova visible in 1054 AD]]
When the core of a massive star collapses, it will form a [[neutron star]], or in the case of cores that exceed the [[Tolman–Oppenheimer–Volkoff limit]], a [[black hole]].  Through a process that is not completely understood, some of the [[gravitational potential energy]] released by this core collapse is converted into a Type Ib, Type Ic, or Type II [[supernova]]. It is known that the core collapse produces a massive surge of [[neutrino]]s, as observed with supernova [[SN 1987A]]. The extremely energetic [[neutrinos]] fragment some nuclei; some of their energy is consumed in releasing [[nucleons]], including [[neutrons]], and some of their energy is transformed into heat and [[kinetic energy]], thus augmenting the [[shock wave]] started by rebound of some of the infalling material from the collapse of the core. Electron capture in very dense parts of the infalling matter may produce additional neutrons. Because some of the rebounding matter is bombarded by the neutrons, some of its nuclei capture them, creating a spectrum of heavier-than-iron material including the radioactive elements up to (and likely beyond) [[uranium]].<ref>[http://www.mpa-garching.mpg.de/HIGHLIGHT/2001/highlight0102_e.html How do Massive Stars Explode?<!-- Bot generated title -->] {{webarchive|url=https://web.archive.org/web/20030627124651/http://www.mpa-garching.mpg.de/HIGHLIGHT/2001/highlight0102_e.html |date=2003-06-27 }}</ref> Although non-exploding red giants can produce significant quantities of elements heavier than iron using neutrons released in side reactions of earlier [[nuclear reactions]], the abundance of elements heavier than [[iron]] (and in particular, of certain isotopes of elements that have multiple stable or long-lived isotopes) produced in such reactions is quite different from that produced in a supernova. Neither abundance alone matches that found in the [[Solar System]], so both supernovae, [[neutron star merger]]s<ref>{{cite web |last=Stromberg |first=Joseph |date=16 July 2013 |title=All the Gold in the Universe Could Come from the Collisions of Neutron Stars |url=http://www.smithsonianmag.com/science-nature/all-the-gold-in-the-universe-could-come-from-the-collisions-of-neutron-stars-13474145/?page=1 |work=[[Smithsonian (magazine)|Smithsonian]] |access-date=27 April 2014}}</ref> and ejection of elements from red giants are required to explain the observed abundance of heavy elements and [[isotopes]] thereof.

The energy transferred from collapse of the core to rebounding material not only generates heavy elements, but provides for their acceleration well beyond [[escape velocity]], thus causing a Type Ib, Type Ic, or Type II supernova. Current understanding of this energy transfer is still not satisfactory; although current computer models of Type Ib, Type Ic, and Type II supernovae account for part of the energy transfer, they are not able to account for enough energy transfer to produce the observed ejection of material.<ref>{{cite web|url=http://www.mpa-garching.mpg.de/HIGHLIGHT/2003/highlight0306_e.html|title=Supernova Simulations Still Defy Explosions|date=June 2003|author=Robert Buras|display-authors=etal|publisher=Max-Planck-Institut für Astrophysik|work=Research Highlights|url-status=dead|archive-url=https://web.archive.org/web/20030803015427/http://www.mpa-garching.mpg.de/HIGHLIGHT/2003/highlight0306_e.html|archive-date=2003-08-03}}</ref> However, neutrino oscillations may play an important role in the energy transfer problem as they not only affect the energy available in a particular flavour of neutrinos but also through other general-relativistic effects on neutrinos.<ref>{{cite journal|doi=10.1023/B:GERG.0000038633.96716.04|title=Addendum to: Gen. Rel. Grav. 28 (1996) 1161, First Prize Essay for 1996: Neutrino Oscillations and Supernovae|journal=General Relativity and Gravitation|volume=36|issue=9|pages=2183–2187|year=2004|last1=Ahluwalia-Khalilova|first1=D. V|bibcode=2004GReGr..36.2183A|arxiv=astro-ph/0404055|s2cid=1045277}}</ref><ref>{{cite journal|bibcode=2017PhRvD..96b3009Y|arxiv=1705.09723|title=GR effects in supernova neutrino flavor transformations|journal=Physical Review D|volume=96|issue=2|pages=023009|last1=Yang|first1=Yue|last2=Kneller|first2=James P|year=2017|doi=10.1103/PhysRevD.96.023009|s2cid=119190550}}
</ref>

Some evidence gained from analysis of the mass and orbital parameters of binary neutron stars (which require two such supernovae) hints that the collapse of an oxygen-neon-magnesium core may produce a supernova that differs observably (in ways other than size) from a supernova produced by the collapse of an iron core.<ref>{{cite journal | author=E. P. J. van den Heuvel | title=X-Ray Binaries and Their Descendants: Binary Radio Pulsars; Evidence for Three Classes of Neutron Stars? | journal=Proceedings of the 5th INTEGRAL Workshop on the INTEGRAL Universe (ESA SP-552) | volume=552 | date=2004 | pages=185–194 | bibcode=2004ESASP.552..185V |arxiv = astro-ph/0407451}}</ref>

The most massive stars that exist today may be completely destroyed by a supernova with an energy greatly exceeding its [[gravitational binding energy]]. This rare event, caused by [[pair-instability supernova|pair-instability]], leaves behind no black hole remnant.<ref name="Hammer">[http://www.mpa-garching.mpg.de/~hammer/lager/pair.pdf Pair Instability Supernovae and Hypernovae.], Nicolay J. Hammer, (2003), accessed May 7, 2007. {{webarchive |url=https://web.archive.org/web/20120608135141/http://www.mpa-garching.mpg.de/~hammer/lager/pair.pdf |date=June 8, 2012 }}</ref> In the past history of the universe, some stars were even larger than the largest that exists today, and they would immediately collapse into a black hole at the end of their lives, due to [[photodisintegration]].