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Schematic illustrating the r-process as it occurs in supernovae or neutron star collisions.[1] Neutrons are rapidly absorbed faster than the resulting nuclei can beta-decay; this allows the r-process to produce very neutron-rich nuclei following the neutron drip line. There are waiting points located at magic numbers N = 50, 82, 126, where beta-decay is favored due to low neutron-capture cross sections resulting from the closed shells; the cycle then repeats until the next waiting point, creating yet heavier nuclei of elements up to the actinides; the natural abundance of these elements results entirely from the r-process. In the superheavy mass region (A ≥ 270), neutron-induced fission or spontaneous fission are expected to become dominant and end the r-process. Most nuclei along the hypothetical r-process path, however, are unknown.[2]

The rapid neutron-capture process, or so-called r-process, is a set of nuclear reactions that in nuclear astrophysics is responsible for the creation of approximately half of the atomic nuclei heavier than iron; the "heavy elements", with the other half produced by the p-process and s-process. The r-process usually synthesizes the most neutron-rich stable isotopes of each heavy element; the r-process can typically synthesize the heaviest four isotopes of every heavy element, and the two heaviest isotopes, which are referred to as r-only nuclei, can only be created via the r-process. Abundance peaks for the r-process occur near atomic weights A = 82 (elements Se, Br and Kr), A = 130 (elements Te, I, and Xe) and A = 196 (elements Os, Ir and Pt).

The r-process entails a succession of rapid neutron captures (hence the name) by one or more heavy seed nuclei, typically beginning with nuclei in the abundance peak centered on 56Fe. The captures must be rapid in the sense that the nuclei must not have time to undergo radioactive decay (typically via β- decay) before another neutron arrives to be captured; this sequence can continue up to the limit of stability of the increasingly neutron-rich nuclei (the neutron drip line) to physically retain neutrons as governed by the short range nuclear force. The r-process therefore must occur in locations where there exist a high density of free neutrons. Early studies theorized that 1024 free neutrons per cm3 would be required, for temperatures about 1GK, in order to match the waiting points, at which no more neutrons can be captured, with the atomic numbers of the abundance peaks for r-process nuclei;[3] this amounts to almost a gram of free neutrons in every cubic centimeter, an astonishing number requiring extreme locations.[a] Traditionally this suggested the material ejected from the reexpanded core of a core-collapse supernova, as part of supernova nucleosynthesis,[4] or decompression of neutron-star matter thrown off by a binary neutron star merger;[5] the relative contributions of these sources to the astrophysical abundance of r-process elements is a matter of ongoing research.[6]

A limited r-process-like series of neutron captures occurs to a minor extent in thermonuclear weapon explosions; these led to the discovery of the elements einsteinium (element 99) and fermium (element 100) in nuclear weapon fallout.



Following pioneering research into the Big Bang and the formation of helium in stars, an unknown process responsible for producing heavier elements found on Earth from hydrogen and helium was suspected to exist. One early attempt at explanation came from Chandrasekhar and Henrich who postulated that elements were produced at temperatures between 6×109 and 8×109 K, their theory accounted for elements up to chlorine, though there was no explanation for elements of atomic weight heavier than 40 amu at non-negligible abundances.[7] This became the foundation of a study by Fred Hoyle, who hypothesized that conditions in the core of collapsing stars would enable nucleosynthesis of the remainder of the elements via rapid capture of densely packed free neutrons. However, there remained unanswered questions about equilibrium in stars that was required to balance beta-decays and precisely account for abundances of elements that would be formed in such conditions.[7]

The need for a physical setting providing rapid neutron capture, which was known to almost certainly have a role in element formation, was also seen in a table of abundances of isotopes of heavy elements by Hans Suess and Harold Urey in 1956,[8] their abundance table revealed larger than average abundances of natural isotopes containing magic numbers[b] of neutrons as well as abundance peaks about 10 amu lighter than stable nuclei containing magic numbers of neutrons which were also in abundance, suggesting that radioactive neutron-rich nuclei having the magic neutron numbers but roughly ten fewer protons were formed. These observations also implied that rapid neutron capture occurred faster than beta decay, and the resulting abundance peaks were caused by so-called waiting points at magic numbers;[3][c] this process, rapid neutron capture by neutron-rich isotopes, became known as the r-process, whereas the s-process was named for its characteristic slow neutron capture. A table apportioning the heavy isotopes phenomenologically between s-process and r-process isotopes was published in 1957 in the B2FH review paper,[3]  which named the r-process and outlined the physics that guides it. Alastair G. W. Cameron also published a smaller study about the r-process in the same year.[9]






Nuclear physics[edit]


Immediately after the severe compression of electrons in a Type II supernova, beta-minus decay is blocked; this is because the high electron density fills all available free electron states up to a Fermi energy which is greater than the energy of nuclear beta decay. However, nuclear capture of those free electrons still occurs, and causes increasing neutronization of matter; this results in an extremely high density of free neutrons which cannot decay, on the order of 1024 neutrons per cm3)[3], and high temperatures. As this re-expands and cools, neutron capture by still-existing heavy nuclei occurs much faster than beta-minus decay; as a consequence, the r-process runs up along the neutron drip line and highly-unstable neutron-rich nuclei are created.

Three processes which affect the climbing of the neutron drip line are a notable decrease in the neutron-capture cross section in nuclei with closed neutron shells, the inhibiting process of photodisintegration, and the degree of nuclear stability in the heavy-isotope region. Neutron captures in r-process nucleosynthesis leads to the formation of neutron-rich, weakly bound nuclei with neutron separation energies as low as 2 MeV.[16][3] At this stage, closed neutron shells at N = 50, 82, and 126 are reached, and neutron capture is temporarily paused; these so-called waiting points are characterized by increased binding energy relative to heavier isotopes, leading to low neutron capture cross sections and a buildup of semi-magic nuclei that are more stable toward beta decay.[17] In addition, nuclei beyond the shell closures are susceptible to quicker beta decay owing to their proximity to the drip line; for these nuclei, beta decay occurs before further neutron capture.[18] Waiting point nuclei are then allowed to beta decay toward stability before further neutron capture can occur,[3] resulting in a slowdown or freeze-out of the reaction.[17]

Decreasing nuclear stability terminates the r-process when its heaviest nuclei become unstable to spontaneous fission, when the total number of nucleons approaches 270; the fission barrier may be low enough before 270 such that neutron capture might induce fission instead of continuing up the neutron drip line.[19] After the neutron flux decreases, these highly unstable radioactive nuclei undergo a rapid succession of beta decays until they reach more stable, neutron-rich nuclei.[20] While the s-process creates an abundance of stable nuclei having closed neutron shells, The r-process, in neutron-rich predecessor nuclei, creates an abundance of radioactive nuclei about 10 amu below the s-process peaks after their decay back to stability.[21]


Astrophysical sites[edit]

The most probable candidate site for the r-process has long been suggested to be core-collapse supernovae (spectral types Ib, Ic and II), which may provide the necessary physical conditions for the r-process. However, the very low abundance of r-process nuclei in the interstellar gas limits the amount each can have ejected, it requires either that only a small fraction of supernovae eject r-process nuclei to the interstellar medium, or that each supernova ejects only a very small amount of r-process material. The ejected material must be relatively neutron-rich, a condition which has been difficult to achieve in models,[4] so that astrophysicists remain uneasy about their adequacy for successful r-process yields.

In 2017, entirely new astronomical data about the r-process was discovered in data about the merger of two neutron stars. Using the gravitational wave data captured in GW170817 to identify the location of the merger, several teams[23][1][24] observed and studied optical data 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, and californium). When released from the huge internal pressure of the neutron star, these ejecta expand and form seed heavy nuclei that rapidly capture free neutrons, and radiate 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.[10] Because of these spectroscopic features it has been argued that such nucleosynthesis in the Milky Way has been primarily ejecta from neutron-star mergers rather than from supernovae.[5]



  2. ^ Neutron number 50, 82 and 126


  1. ^ a b Pian, E.; et al. (2017). "Spectroscopic identification of r-process nucleosynthesis in a double neutron-star merger" (PDF). Nature. 551 (7678): 67–70. arXiv:1710.05858. Bibcode:2017Natur.551...67P. doi:10.1038/nature24298. PMID 29094694.
  2. ^ Thoennessen, M. (2004). "Reaching the limits of nuclear stability" (PDF). Reports on Progress in Physics. 67: 1187–1232. doi:10.1088/0034-4885/67/7/R04.
  5. ^ a b Kasen, D.; Metzger, B.; Barnes, J.; Quataert, E.; Ramirez-Ruiz, E. (2017). "Origin of the heavy elements in binary neutron-star mergers from a gravitational-wave event" (PDF). Nature. 551 (7678): 80–84. arXiv:1710.05463. Bibcode:2017Natur.551...80K. doi:10.1038/nature24453. PMID 29094687.
  7. ^ a b Hoyle, F. (1946). "The Synthesis of the Elements from Hydrogen" (pdf). Monthly Notices of the Royal Astronomical Society. 106 (5): 343–383. doi:10.1093/mnras/106.5.343.
  12. ^ See Table 4 in Seeger, Fowler & Clayton 1965.
  15. ^ Bartlett, A.; Görres, J.; Mathews, G.J.; Otsuki, K.; Wiescher, W. (2006). "Two-neutron capture reactions and the r process" (PDF). Physical Review C. 74: 015082. doi:10.1103/PhysRevC.74.015802.
  18. ^ Wang, R.; Chen, L.W. (2015). "Positioning the neutron drip line and the r-process paths in the nuclear landscape". Physical Review C. 92 (3): 031303–1—031303–5. arXiv:1410.2498. doi:10.1103/PhysRevC.92.031303.
  23. ^ Arcavi, I.; et al. (2017). "Optical emission from a kilonova following a gravitational-wave-detected neutron-star merger" (PDF). Nature. 551 (7678): 64–66. arXiv:1710.05843. Bibcode:2017Natur.551...64A. doi:10.1038/nature24291.
  24. ^ Smartt, S. J.; et al. (2017). "A kilonova as the electromagnetic counterpart to a gravitational-wave source" (PDF). Nature. 551 (7678): 75–79. arXiv:1710.05841. Bibcode:2017Natur.551...75S. doi:10.1038/nature24303. PMID 29094693.
  28. ^ Tanvir, N.; et al. (2013). "A 'kilonova' associated with the short-duration gamma-ray burst GRB 130603B" (PDF). Nature. 500 (7464): 547–9. arXiv:1306.4971. Bibcode:2013Natur.500..547T. doi:10.1038/nature12505. PMID 23912055.