Supernova nucleosynthesis

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Supernova nucleosynthesis is a theory of the production of many different chemical elements in supernova explosions, first advanced by Fred Hoyle in 1954.[1] The nucleosynthesis, or fusion of lighter elements into heavier ones, occurs during explosive oxygen burning and silicon burning.[2] Those fusion reactions create the elements silicon, sulfur, chlorine, argon, sodium, potassium, calcium, scandium, titanium and iron peak elements: vanadium, chromium, manganese, iron, cobalt, and nickel. These are called "primary elements", in that they can be fused from pure hydrogen and helium in massive stars. As a result of their ejection from supernovae, their abundances increase within the interstellar medium. Elements heavier than nickel are created primarily by a rapid capture of neutrons in a process called the r-process. However, these are much less abundant than the primary chemical elements. Other processes thought to be responsible for some of the nucleosynthesis of underabundant heavy elements, notably a proton capture process known as the rp-process and a photodisintegration process known as the gamma (or p) process. The latter synthesizes the lightest, most neutron-poor, isotopes of the heavy elements.


Main article: Supernova

A supernova is a massive explosion of a star that occurs under two principal scenarios. The first is that a white dwarf star undergoes a nuclear-based explosion after it reaches its Chandrasekhar limit after absorbing mass from a neighboring star (usually a red giant). The second, and more common, cause is when a massive star, usually a supergiant, reaches nickel-56 in its nuclear fusion (or burning) processes. This isotope undergoes radioactive decay into iron-56, which has one of the highest binding energies of all of the isotopes, and is the last element that produces a net release of energy by nuclear fusion, exothermically.

All nuclear fusion reactions that produce heavier elements cause the star to lose energy and are said to be endothermic reactions. The pressure that supports the star's outer layers drops sharply. As the outer envelope is no longer sufficiently supported by the radiation pressure, the star's gravity pulls its outer layers rapidly inward. As the star collapses, these outer layers collide with the incompressible stellar core, producing a shockwave that expands outward through the unfused material of the outer shell. The pressures and densities in the shockwave are sufficient to induce fusion in that material, and the energy released leads to the star's explosion, dispersing material from the star into interstellar space.

Silicon photodisintegration rearrangment and quasiequilibrium[edit]

After a star completes the oxygen burning process, its core is composed primarily of silicon and sulfur.[3] If it has sufficiently high mass, it further contracts until its core reaches temperatures in the range of 2.7–3.5 GK (230–300 keV). At these temperatures, silicon and other elements photodisintegrate by energetic thermal photons ejecting alpha particles.[3] Silicon burning differs from earlier fusion stages of nucleosynthesis in that it entails a balance between alpha-particle captures and their inverse photo ejection which establishes abundances all alpha-particle elements in the following sequence in which each alpha particle capture shown is opposed by its inverse reaction, namely, photo ejection of an alpha particle by abundant thermal photons:

28Si + 4He ↔ 32S + photon; 32S + 4He ↔ 36Ar + photon; 36Ar + 4He ↔ 40Ca + photon; 40Ca + 4He ↔ 44Ti + photon; 44Ti + 4He ↔ 48Cr + photon; 48Cr + 4He ↔ 52Fe + photon; 52Fe + 4He ↔ 56Ni + photon; 56Ni + 4He ↔ 60Zn + photon

In these circumstances of rapid opposing reactions 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 state is called "quasiequilibrium".[4] The quasiequilibrium buildup shuts off at 56Ni because the alpha-particle captures become slower whereas the photo ejection from heavier nuclei becomes faster. Non-alpha nuclei are also involved via many reactions similar to 36Ar + neutron ↔ 37Ar + photon and its inverse, where the free densities of protons and neutrons are also set by the quasiequilibrium. The silicon-burning quasiequilibrium is a unique construction, simultaneously the most abstract and the most beautiful of nucleosynthesis processes.

The entire silicon-burning sequence lasts about one day in the core of a contracting massive star and stops after nickel-56 has become the dominant abundance. The explosive burning caused when the supernova shock passes through the silicon-burning shell lasts only seconds but is the major contributor to nucleosynthesis in the mass range 28-60.[5] The star can no longer release energy via nuclear fusion because a nucleus with 56 nucleons has the lowest mass per nucleon of all the elements in the sequence. The next step up in the alpha-particle chain would be 60Zn, which has slightly more mass per nucleon and thus is less thermodynamically favorable. 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.[6] See SN 1978A light curve for the aftermath of that opportunity.

During this phase of the core contraction, the potential energy of gravitational contraction heats the interior to 5 GK (430 keV) and this opposes and delays the 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. The central portion of the star is now crushed into either a neutron star or, if the star is massive enough, 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 that lasts days to months. The supernova explosion releases a large burst of neutrons, which synthesizes, in about one second while-inside the star, roughly half of the supply of elements in the universe that are heavier than iron via a neutron-capture mechanism known as the r-process (where the “r” stands for rapid neutron capture).

Nuclides Synthesized[edit]

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. Nucleosynthesis within those lighter stars is therefore limited to nuclides that are fused in material located above the final white dwarf. This limits their yields to interstellar gas to carbon, nitrogen, and to isotopes heavier than iron by slow capture of neutrons (the s process). Virtually all of stellar nucleosynthesis occurs in stars that are massive enough to end in Type II supernovae.[7][8] In the pre supernova massive star this includes carbon burning, oxygen burning and silicon burning. Much of that yield may never leave the star but disappear into its collapsed core. The yield that is ejected is substantially fused also in explosive burning caused by the shock wave launched by core collapse[9] (see Supernova). Prior to core collapse, fusion of elements between silicon and iron occurs only in the largest of stars, and then in limited amounts. Thus the nucleosynthesis of the abundant and primary elements,[10] defined as those that could be synthesized in stars of only hydrogen and helium (left by the Big Bang), is substantially limited to Supernova nucleosynthesis, as Fred Hoyle first described[11] in his pioneering work establishing this subject.

The r-process[edit]

Main article: r-process
A version of the periodic table indicating the origins – including supernova nucleosynthesis – of the elements. All elements above 103 (lawrencium) are also manmade and are not included.

During supernova nucleosynthesis, the r-process (r for rapid) 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 flux is astonishingly high, about 1022 neutrons per square centimeter per second. First calculation of a dynamic r-process, showing the evolution of calculated results with time,[12] also suggested that the r-process abundances are a superposition of differing neutron fluences. Small fluence produces the first r-process abundance peak near atomic weight A=130 but no actinides, whereas large fluence produces the actinides uranium and thorium but no longer contains the A=130 abundance peak. These processes occur in a fraction of a second to a few seconds, depending on details. Hundreds of subsequent papers published have utilized this time-dependent approach. Interestingly, the only modern nearby supernova, 1987A, has not revealed r-process enrichments. Modern thinking is that the r-process yield may be ejected from some supernovae but swallowed up in others as part of the residual neutron star or black hole.

See also[edit]



  1. ^ "Synthesis of the laments from carbon to nickel" Astrophys. J. Suppl. 1, 121 (1954)
  2. ^ Woosley, S.E.; W. D. Arnett & D. D. Clayton (1973). "Explosive burning of oxygen and silicon". The Astrophysical Journal Supplement 26: 231–312. Bibcode:1973ApJS...26..231W. doi:10.1086/190282. 
  3. ^ a b Donald D. Clayton "Principles of Stellar Evolution and Nucleosynthesis" Univ. of Chicago Press (Chicago: 1983) Chap. 7, p.519-524
  4. ^ "Nucleosynthesis During Silicon Burning", D. Bodansky. D.D, Clayton & W.A. Fowler, Phys. Rev. Letters, 20, 161, (1968); “Nuclear quasi-equilibrium during silicon burning”, D. Bodansky. D.D, Clayton & W.A. Fowler, Astrophys. J. Suppl. No. 148, 16, 299, (1968); Chapter 7 of Clayton's 1968 textbook, Principles of Stellar Evolution and Nucleosynthesis
  5. ^ "Nucleosynthesis During Silicon Burning", D. Bodansky. D.D, Clayton & W.A. Fowler, Phys. Rev. Letters, 20, 161, (1968); “Nuclear quasi-equilibrium during silicon burning”, D. Bodansky. D.D, Clayton & W.A. Fowler, Astrophys. J. Suppl. No. 148, 16, 299, (1968); "Explosive Burning of Oxygen and Silicon"Woosley S.E. Arnett W.D. Clayton D.D. Astrophys. J. Suppl. 26, 231 (1973); "Handbook of Isotopes in the Cosmos" Donald D. Clayton (Cambridge University Press, Cambridge 2003)
  6. ^ Clayton, Colgate & Fishman "Gamma Ray Lines from young Supernova Remnants", Astrophys. J. 155, 75 (1969)
  7. ^ Donald D. Clayton, Handbook of Isotopes in the Cosmos, Cambridge University Press (2003)
  8. ^ François, P.; et al. (2004). "The evolution of the Milky Way from its earliest phases: Constraints on stellar nucleosynthesis". Astronomy and Astrophysics 421 (2): 613–621. arXiv:astro-ph/0401499. Bibcode:2004A&A...421..613F. doi:10.1051/0004-6361:20034140
  9. ^ Woosley SE, Arnett D, Clayton DD, "Explosive burning go oxygen and silicon" Astrophys. J. Suppl. 26, 231 (1973)
  10. ^ Donald D. Clayton, "Fred Hoyle, primary nucleosynthesis and radioactivity", New Astronomy Reviews52, 360-63 (2008)
  11. ^ F. Hoyle, "Synthesis of elements from carbon to nickel", Astrophys. J. Suppl. 1, 121 (1954)
  12. ^ P. A. Seeger; W.A. Fowler; D. D. Clayton (1965). "Nucleosynthesis of heavy elements by neutron capture". The Astrophysical Journal Supplement 11: 121–166. Bibcode:1965ApJS...11..121S. doi:10.1086/190111. 

Other reading[edit]

  • E. M. Burbidge, G. R. Burbidge, W. A. Fowler, F. Hoyle, Synthesis of the Elements in Stars, Rev. Mod. Phys. 29 (1957) 547 (article at the Physical Review Online Archive).
  • D. D. Clayton, "Handbook of Isotopes in the Cosmos", Cambridge University Press, 2003, ISBN 0-521-82381-1.

External links[edit]