Supernova nucleosynthesis is a theory of the production of many different chemical elements in supernova explosions, first advanced by Fred Hoyle in 1954. The nucleosynthesis, or fusion of lighter elements into heavier ones, occurs during explosive oxygen burning and silicon burning. 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.
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.
Nuclear fusion sequence and the alpha process
After a star completes the oxygen burning process, its core is composed primarily of silicon and sulfur. 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 can photodisintegrate, emitting a proton or alpha particle. Silicon burning entails the alpha process, which creates new elements by adding one of these alpha particles (the equivalent of a helium nucleus, two protons plus two neutrons) per step in the following sequence:
The entire silicon-burning sequence lasts about one day and stops when nickel-56 has been produced. The star can no longer release energy via nuclear fusion because a nucleus with 56 nucleons has the lowest mass per nucleon (any proton or neutron) of all the elements in the alpha process sequence. Although iron-58 and nickel-62 have slightly higher binding energies per nucleon than iron-56, the next step up in the alpha process would be zinc-60, which has slightly more mass per nucleon and thus, is less thermodynamically favorable. Nickel-56 (which has 28 protons) has a half-life of 6.02 days and decays via β+ decay to cobalt-56 (27 protons), which in turn has a half-life of 77.3 days as it decays to iron-56 (26 protons). However, only minutes are available for the nickel-56 to decay within the core of a massive star. The star has run out of nuclear fuel and within minutes begins to contract.
During this phase of the 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 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).
The maximum weight for an element produced by fusion in a normal star is that of iron, reaching an isotope with an atomic mass of 56 (see Stellar nucleosynthesis). Prior to a supernova, fusion of elements between silicon and iron occurs only in the largest of stars, in the silicon burning process. (A slow neutron capture process, known as the s-process which also occurs during normal stellar nucleosynthesis can create elements up to bismuth with an atomic mass of approximately 209. However, the s-process occurs primarily in low-mass stars that evolve more slowly.) Once the core fails to produce enough energy to support the outer envelope of gases the star explodes as a supernova producing the bulk of elements beyond iron. Production of elements from iron to uranium occurs within seconds in a supernova explosion. Due to the large amounts of energy released, much higher temperatures and densities are reached than at normal stellar temperatures. These conditions allow for an environment where transuranium elements might be formed.
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 weight. 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, 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.
- Energy is produced in the isolated fusion reaction of nickel-56 with helium-4, but production of the latter (by photodisintegration of heavier nuclei) is costly, and consumes energy, causing alpha buildup of nickel to be shut off due to the essential fact that nickel-56 has nucleon binding energy less zinc-60.
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