Supernova nucleosynthesis
Supernova nucleosynthesis is the nucleosynthesis of chemical elements in supernova explosions.
In sufficiently massive stars, the nucleosynthesis by fusion of lighter elements into heavier ones occurs during sequential
During hydrostatic burning these fuels synthesize overwhelmingly the alpha nuclides (A = 2Z), nuclei composed of integer numbers of helium-4 nuclei. Initially, two helium-4 nuclei fuse into a single beryllium-8 nucleus. The addition of another helium 4 nucleus to the beryllium yields carbon-12, followed by oxygen-16, neon-20 and so on, each time adding 2 protons and 2 neutrons to the growing nucleus. A rapid final explosive burning[1] is caused by the sudden temperature spike owing to passage of the radially moving shock wave that was launched by the gravitational collapse of the core. W. D. Arnett and his Rice University colleagues[2][1] demonstrated that the final shock burning would synthesize the non-alpha-nucleus isotopes more effectively than hydrostatic burning was able to do,[3][4] suggesting that the expected shock-wave nucleosynthesis is an essential component of supernova nucleosynthesis. Together, shock-wave nucleosynthesis and hydrostatic-burning processes create most of the isotopes of the elements carbon (Z = 6), oxygen (Z = 8), and elements with Z = 10 to 28 (from neon to nickel).[4][5] As a result of the ejection of the newly synthesized isotopes of the chemical elements by supernova explosions, their abundances steadily increased within interstellar gas. That increase became evident to astronomers from the initial abundances in newly born stars exceeding those in earlier-born stars.
Elements heavier than nickel are comparatively rare owing to the decline with atomic weight of their nuclear binding energies per nucleon, but they too are created in part within supernovae. Of greatest interest historically has been their synthesis by rapid capture of
History
In 1946, Fred Hoyle proposed that elements heavier than hydrogen and helium would be produced by nucleosynthesis in the cores of massive stars.[6] It had previously been thought that the elements we see in the modern universe had been largely produced during its formation. At this time, the nature of supernovae was unclear and Hoyle suggested that these heavy elements were distributed into space by rotational instability. In 1954, the theory of nucleosynthesis of heavy elements in massive stars was refined and combined with more understanding of supernovae to calculate the abundances of the elements from carbon to nickel.[7] Key elements of the theory included:
- the prediction of the excited state in the 12C nucleus that enables the triple-alpha process to burn resonantly to carbon and oxygen;
- the thermonuclear sequels of carbon-burningsynthesizing Ne, Mg and Na; and
- oxygen-burning synthesizing silicon, aluminum, and sulphur.
The theory predicted that
In 1957, a paper by the authors
Thirteen years after the B²FH paper, W.D. Arnett and colleagues[2][1] demonstrated that the final burning in the passing shock wave launched by collapse of the core could synthesize non-alpha-particle isotopes more effectively than hydrostatic burning could,[3][4] suggesting that explosive nucleosynthesis is an essential component of supernova nucleosynthesis. A shock wave rebounded from matter collapsing onto the dense core, if strong enough to lead to mass ejection of the mantle of supernovae, would necessarily be strong enough to provide the sudden heating of the shells of massive stars needed for explosive thermonuclear burning within the mantle. Understanding how that shock wave can reach the mantle in the face of continuing infall onto the shock became the theoretical difficulty. Supernova observations assured that it must occur.
White dwarfs were proposed as possible progenitors of certain supernovae in the late 1960s,[12] although a good understanding of the mechanism and nucleosynthesis involved did not develop until the 1980s.[13] This showed that type Ia supernovae ejected very large amounts of radioactive nickel and lesser amounts of other iron-peak elements, with the nickel decaying rapidly to cobalt and then iron.[14]
Era of computer models
The papers of Hoyle (1946) and Hoyle (1954) and of B²FH (1957) were written by those scientists before the advent of the age of computers. They relied on hand calculations, deep thought, physical intuition, and familiarity with details of nuclear physics. Brilliant as these founding papers were, a cultural disconnect soon emerged with a younger generation of scientists who began to construct computer programs[15] that would eventually yield numerical answers for the advanced evolution of stars[16] and the nucleosynthesis within them.[17][18]
Cause
A supernova is a violent explosion of a star that occurs under two principal scenarios. The first is that a
The nickel-56 isotope has one of the largest
Nuclear fusion reactions that produce elements heavier than iron absorb nuclear energy and are said to be
Silicon burning
After a star completes the
28Si + 4He ⇌ 32S + γ 32S + 4He ⇌ 36Ar + γ 36Ar + 4He ⇌ 40Ca + γ 40Ca + 4He ⇌ 44Ti + γ 44Ti + 4He ⇌ 48Cr + γ 48Cr + 4He ⇌ 52Fe + γ 52Fe + 4He ⇌ 56Ni + γ 56Ni + 4He ⇌ 60Zn + γ
The alpha-particle nuclei 44Ti and those more massive in the final five reactions listed are all radioactive, but they decay after their ejection in supernova explosions into abundant isotopes of Ca, Ti, Cr, Fe and Ni. This post-supernova radioactivity became of great importance for the emergence of gamma-ray-line astronomy.[22]
In these physical circumstances of rapid opposing reactions, namely alpha-particle capture and photo ejection of alpha particles, the abundances are not determined by alpha-particle-capture cross sections; rather they are determined by the values that the abundances must assume in order to balance the speeds of the rapid opposing-reaction currents. Each abundance takes on a stationary value that achieves that balance. This picture is called nuclear quasiequilibrium.[23][24][25] Many computer calculations, for example,[26] using the numerical rates of each reaction and of their reverse reactions have demonstrated that quasiequilibrium is not exact but does characterize well the computed abundances. Thus, the quasiequilibrium picture presents a comprehensible picture of what actually happens. It also fills in an uncertainty in Hoyle's 1954 theory. The quasiequilibrium buildup shuts off after 56Ni because the alpha-particle captures become slower whereas the photo ejections from heavier nuclei become faster. Non-alpha-particle nuclei also participate, using a host of reactions similar to
- 36Ar + neutron ⇌ 37Ar + photon
and its inverse which set the stationary abundances of the non-alpha-particle isotopes, where the free densities of protons and neutrons are also established by the quasiequilibrium. However, the abundance of free neutrons is also proportional to the excess of neutrons over protons in the composition of the massive star; therefore the abundance of 37Ar, using it as an example, is greater in ejecta from recent massive stars than it was from those in early stars of only H and He; therefore 37Cl, to which 37Ar decays after the nucleosynthesis, is called a "secondary isotope".
In interest of brevity, the next stage, an intricate photo-disintegration rearrangement, and the nuclear quasiequilibrium that it achieves, are referred to as silicon burning. The silicon burning in the star progresses through a temporal sequence of such nuclear quasiequilibria in which the abundance of 28Si slowly declines and that of 56Ni slowly increases. This amounts to a nuclear abundance change 2 28Si ≫ 56Ni, which may be thought of as silicon burning into nickel ("burning" in the nuclear sense). The entire silicon-burning sequence lasts about one day in the core of a contracting massive star and stops after 56Ni has become the dominant abundance. The final explosive burning caused when the supernova shock passes through the silicon-burning shell lasts only seconds, but its roughly 50% increase in the temperature causes furious nuclear burning, which becomes the major contributor to nucleosynthesis in the mass range 28–60
After the final 56Ni stage, the star can no longer release energy via nuclear fusion, because a nucleus with 56 nucleons has the lowest
56Ni (which has 28 protons) has a half-life of 6.02 days and decays via β+ decay to 56Co (27 protons), which in turn has a half-life of 77.3 days as it decays to 56Fe (26 protons). However, only minutes are available for the 56Ni to decay within the core of a massive star.
This establishes 56Ni as the most abundant of the radioactive nuclei created in this way. Its radioactivity energizes the late supernova light curve and creates the pathbreaking opportunity for gamma-ray-line astronomy.[22] See SN 1987A light curve for the aftermath of that opportunity.
Clayton and Meyer[26] have recently generalized this process still further by what they have named the secondary supernova machine, attributing the increasing radioactivity that energizes late supernova displays to the storage of increasing Coulomb energy within the quasiequilibrium nuclei called out above as the quasiequilibria shift from primarily 28Si to primarily 56Ni. The visible displays are powered by the decay of that excess Coulomb energy.
During this phase of the core contraction, the potential energy of gravitational compression heats the interior to roughly three billion kelvins, which briefly maintains pressure support and opposes rapid core contraction. However, since no additional heat energy can be generated via new fusion reactions, the final unopposed contraction rapidly accelerates into a collapse lasting only a few seconds. At that point, the central portion of the star is crushed into either a neutron star or, if the star is massive enough, into a black hole.
The outer layers of the star are blown off in an explosion triggered by the outward moving supernova shock, known as a Type II supernova whose displays last days to months. The escaping portion of the supernova core may initially contain a large density of free neutrons, which may synthesize, in about one second while inside the star, roughly half of the elements in the universe that are heavier than iron via a rapid neutron-capture mechanism known as the r-process. See below.
Nuclides synthesized
Stars with initial masses less than about eight times the sun never develop a core large enough to collapse and they eventually lose their atmospheres to become white dwarfs, stable cooling spheres of carbon supported by the pressure of
A significant minority of white dwarfs will explode, however, either because they are in a binary orbit with a companion star that loses mass to the stronger gravitational field of the white dwarf, or because of a merger with another white dwarf. The result is a white dwarf which exceeds its Chandrasekhar limit and explodes as a type Ia supernova, synthesizing about a solar mass of radioactive 56Ni isotopes, together with smaller amounts of other iron peak elements. The subsequent radioactive decay of the nickel to iron keeps Type Ia optically very bright for weeks and creates more than half of all the iron in the universe.[28]
Virtually all of the remainder of stellar nucleosynthesis occurs, however, in stars that are massive enough to end as
r-process nucleosynthesis
This section needs additional citations for verification. (October 2022) |
During supernova nucleosynthesis, the r-process creates very neutron-rich heavy isotopes, which decay after the event to the first stable isotope, thereby creating the neutron-rich stable isotopes of all heavy elements. This neutron capture process occurs in high neutron density with high temperature conditions.
In the r-process, any heavy nuclei are bombarded with a large neutron flux to form highly unstable neutron rich nuclei which very rapidly undergo beta decay to form more stable nuclei with higher atomic number and the same atomic mass. The neutron density is extremely high, about 1022–24 neutrons per cubic centimeter.
Initial calculations of an evolving r-process, showing the evolution of calculated results with time,
Entirely new astronomical data about the r-process was discovered in 2017 when the LIGO and Virgo gravitational-wave observatories discovered a merger of two neutron stars that had previously been orbiting one another.[31] That can happen when both massive stars in orbit with one another become core-collapse supernovae, leaving neutron-star remnants.
The localization on the sky of the source of those gravitational waves radiated by that orbital collapse and merger of the two neutron stars, creating a black hole, but with significant spun off mass of highly neutronized matter, enabled several teams[32][33][34] to discover and study the remaining optical counterpart of the merger, finding spectroscopic evidence of r-process material thrown off by the merging neutron stars.
The bulk of this material seems to consist of two types: Hot blue masses of highly radioactive r-process matter of lower-mass-range heavy nuclei (A < 140) and cooler red masses of higher mass-number r-process nuclei (A > 140) rich in actinides (such as uranium, thorium, californium etc.). When released from the huge internal pressure of the neutron star, this neutron-rich ejecta expands and radiates detected optical light for about a week. Such duration of luminosity would not be possible without heating by internal radioactive decay, which is provided by r-process nuclei near their waiting points. Two distinct mass regions (A < 140 and A > 140) for the r-process yields have been known since the first time dependent calculations of the r-process.[30] Because of these spectroscopic features it has been argued that r-process nucleosynthesis in the Milky Way may have been primarily ejecta from neutron-star mergers rather than from supernovae.[35]
See also
References
- ^ a b c d e
Woosley, S.E.; Arnett, W.D.; Clayton, D.D. (1973). "The Explosive burning of oxygen and silicon". S2CID 222372611.
- ^ a b c
Arnett, W. D.; Clayton, D. D. (1970). "Explosive Nucleosynthesis in Stars". S2CID 38865963.
- ^ a b See Figures 1, 3, and 4 in Arnett & Clayton (1970) and Fig. 2, p. 241 in Woosley, Arnett & Clayton 1973
- ^ doi:10.1086/192237. Archivedfrom the original on 2023-01-13. Retrieved 2019-07-11.
- ^
Thielemann, Fr.-K.; Nomoto, K.; Hashimoto, M.-A. (1996). "Core-Collapse Supernovae and Their Ejecta". doi:10.1086/176980.
- ^ a b .
- ^ a b
doi:10.1086/190005.
- ^ .
- ^
S2CID 118423007.
- ^ See Clayton 2008, p. 363, footnote 1
- B²FH paperarticle
- ^
Finzi, A.; Wolf, R.A. (1967). "Type I supernovae". doi:10.1086/149317.
- ^
Nomoto, Ken'Ichi (1980). "White dwarf models for type I supernovae and quiet supernovae, and presupernova evolution". S2CID 120969575.
- ^
Nomoto, K.; Thielemann, F.-K.; Yokoi, K. (1984). "Accreting white dwarf models of type I supernovae. III - Carbon deflagration supernovae". doi:10.1086/162639.
- ISBN 9780226109534. Archivedfrom the original on 2023-01-13. Retrieved 2022-03-30.
- ^ for example
Iben, I. Jr. (1967). "Stellar Evolution. VI. Evolution from the Main Sequence to the Red-Giant Branch for Stars of Mass 1 M☉, 1.25 M☉, and 1.5 M☉*". Astrophysical Journal. 147: 624. doi:10.1086/149040. which contains a description of helium burning.
- doi:10.1086/192237. Archived(PDF) from the original on 2023-01-13. Retrieved 2021-06-24.
- ^
Thielemann, Fr.-K.; Nomoto, K.; Hashimoto, M.-A. (1996). "Core-collapse supernovae and their ejecta". doi:10.1086/176980.
- ^
Heger, A.; Fryer, C.L.; Woosley, S.E.; Langer, N.; Hartmann, D.H. (2003). "How Massive Single Stars End Their Life". S2CID 59065632.
- from the original on 2023-01-13. Retrieved 2018-02-18.
- ^ a b
Clayton, D.D. (1983). Principles of Stellar Evolution and Nucleosynthesis. ISBN 0226109534– via Archive.org.
- ^ doi:10.1086/149849. Archivedfrom the original on 2023-01-13. Retrieved 2018-11-04.
- ^ from the original on 2023-01-13. Retrieved 2019-07-11.
- ^ a b
Bodansky, D.; Clayton, D. D.; Fowler, W. A. (1968). "Nuclear Quasi-Equilibrium during Silicon Burning". doi:10.1086/190176.
- ^ Clayton, D. D. (1968). "Chapter 7". Principles of Stellar Evolution and Nucleosynthesis. University of Chicago Press.
- ^ a b Clayton, D. D.; Meyer, B. S. (2016). "The secondary supernova machine: Gravitational compression, stored Coulomb energy, and SNII displays". .
- ^ a b Clayton, D. D. (2003). Handbook of Isotopes in the Cosmos. Cambridge University Press.
- ^ a b
François, P.; Matteucci, F.; Cayrel, R.; Spite, M.; Spite, F.; Chiappini, C. (2004). "The evolution of the Milky Way fromits earliest phases: Constraints on stellar nucleosynthesis". S2CID 16257700.
- .
- ^ doi:10.1086/190111. Archivedfrom the original on 2021-04-28. Retrieved 2019-07-11.
- ^
Abbott, B. P.; et al. (2017). "GW170817: Observation of Gravitational Waves from a Binary Neutron Star Inspiral". S2CID 217163611.
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Arcavi, I.; et al. (2017). "Optical emission from a kilonova following a gravitational-wave-detected neutron-star merger". S2CID 205261241.
- ^
Pian, E.; et al. (2017). "Spectroscopic identification of r-process nucleosynthesis in a double neutron-star merger". S2CID 3840214.
- ^
Smartt, S. J.; et al. (2017). "A kilonova as the electromagnetic counterpart to a gravitational-wave source". S2CID 205261388.
- from the original on 2023-01-13. Retrieved 2018-09-27.
Other reading
- Burbidge, E. M.; Burbidge, G. R.; Fowler, W.A.; Hoyle, F. (1957). "Synthesis of the Elements in Stars". .
- Clayton, D. (2003). Handbook of isotopes in the cosmos. ISBN 978-0-521-82381-4.
External links
- "Atom smashers shed light on supernovae, Big Bang". Sky & Telescope Online. 22 April 2005.
- Gonzalez, G.; Brownlee, D.; Ward, P. (2001). "The galactic habitable zone: Galactic chemical evolution". S2CID 18179704.