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r-process

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Nuclear physics
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Innuclear astrophysics,therapid neutron-capture process,also known as ther-process,is a set ofnuclear reactionsthat is responsible forthe creationof approximately half of theatomic nucleiheavier than iron,the "heavy elements", with the other half produced by thep-processands-process.Ther-process usually synthesizes the most neutron-rich stable isotopes of each heavy element. Ther-process can typically synthesize the heaviest four isotopes of every heavy element; of these, the heavier two are calledr-only nucleibecause they are created exclusively via ther-process. Abundance peaks for ther-process occur nearmass numbersA= 82(elements Se, Br, and Kr),A= 130(elements Te, I, and Xe) andA= 196(elements Os, Ir, and Pt).

Ther-process entails a succession ofrapidneutron captures(hence the name) by one or more heavyseed nuclei,typically beginning with nuclei in the abundance peak centered on56Fe.The captures must be rapid in the sense that the nuclei must not have time to undergoradioactive decay(typically via β−decay) before anotherneutronarrives to be captured. This sequence can continue up to the limit of stability of the increasingly neutron-rich nuclei (theneutron drip line) to physically retain neutrons as governed by the short range nuclear force. Ther-process therefore must occur in locations where there exists a high density offree neutrons.Early studies theorized that 1024free neutrons per cm3would be required, for temperatures of about 1 GK, in order to match the waiting points, at which no more neutrons can be captured, with the mass numbers of the abundance peaks forr-process nuclei.[1]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 acore-collapse supernova,as part ofsupernova nucleosynthesis,[2]or decompression ofneutron starmatter thrown off by a binaryneutron star mergerin akilonova.[3]The relative contribution of each of these sources to the astrophysical abundance ofr-process elements is a matter of ongoing research as of 2018.[4]

A limitedr-process-like series of neutron captures occurs to a minor extent inthermonuclear weaponexplosions. These led to the discovery of the elementseinsteinium(element 99) andfermium(element 100) innuclear weapon fallout.

Ther-process contrasts with thes-process,the other predominant mechanism for the production of heavy elements, which is nucleosynthesis by means ofslowcaptures of neutrons. In general, isotopes involved in thes-process have half-lives long enough to enable their study in laboratory experiments, but this is not typically true for isotopes involved in ther-process.[5]Thes-process primarily occurs within ordinary stars, particularlyAGB stars,where the neutron flux is sufficient to cause neutron captures to recur every 10–100 years, much too slow for ther-process, which requires 100 captures per second. Thes-process issecondary,meaning that it requires pre-existing heavy isotopes as seed nuclei to be converted into other heavy nuclei by a slow sequence of captures of free neutrons. Ther-process scenarios create their own seed nuclei, so they might proceed in massive stars that contain no heavy seed nuclei. Taken together, ther- ands-processes account for almost the entireabundance of chemical elementsheavier than iron. The historical challenge has been to locate physical settings appropriate to their time scales.

History

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Following pioneering research into theBig Bangand the formation ofheliumin stars, an unknown process responsible for producing heavier elements found on Earth fromhydrogenand helium was suspected to exist. One early attempt at explanation came fromSubrahmanyan Chandrasekharand Louis R. Henrich who postulated that elements were produced at temperatures between 6×109and 8×109K.Their theory accounted for elements up tochlorine,though there was no explanation for elements ofatomic weightheavier than 40amuat non-negligible abundances.[6] This became the foundation of a study byFred 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 forabundances of elementsthat would be formed in such conditions.[6]

The need for a physical setting providing rapidneutron 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 byHans SuessandHarold Ureyin 1956.[7]Their abundance table revealed larger than average abundances of natural isotopes containingmagic numbers[b]of neutrons as well as abundance peaks about 10 amu lighter thanstable nucleicontaining 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 thanbeta decay,and the resulting abundance peaks were caused by so-calledwaiting pointsat magic numbers.[1][c]This process, rapid neutron capture by neutron-rich isotopes, became known as ther-process, whereas thes-process was named for its characteristic slow neutron capture. A table apportioning the heavy isotopes phenomenologically betweens-process andr-process isotopes was published in 1957 in theB2FH review paper,[1]  which named ther-process and outlined the physics that guides it.[8]Alastair G. W. Cameronalso published a smaller study about ther-process in the same year.[9]

The stationaryr-process as described by the B2FH paper was first demonstrated in a time-dependent calculation atCaltechby Phillip A. Seeger,William A. FowlerandDonald D. Clayton,[10]who found that no single temporal snapshot matched the solarr-process abundances, but, that when superposed, did achieve a successful characterization of ther-process abundance distribution. Shorter-time distributions emphasize abundances at atomic weights less thanA= 140,whereas longer-time distributions emphasized those at atomic weights greater thanA= 140.[11]Subsequent treatments of ther-process reinforced those temporal features. Seeger et al. were also able to construct more quantitative apportionment betweens-process andr-process of the abundance table of heavy isotopes, thereby establishing a more reliable abundance curve for ther-process isotopes than B2FH had been able to define. Today, ther-process abundances are determined using their technique of subtracting the more reliables-process isotopic abundances from the total isotopic abundances and attributing the remainder tor-process nucleosynthesis.[12]Thatr-process abundance curve (vs. atomic weight) has provided for many decades the target for theoretical computations of abundances synthesized by the physicalr-process.

The creation of free neutrons by electron capture during the rapid collapse to high density of a supernova core along with quick assembly of some neutron-rich seed nuclei makes ther-process aprimary nucleosynthesis process,a process that can occur even in a star initially of pure H and He. This in contrast to the B2FH designation which is asecondary processbuilding on preexisting iron. Primary stellar nucleosynthesis begins earlier in the galaxy than does secondary nucleosynthesis. Alternatively the high density of neutrons within neutron stars would be available for rapid assembly intor-process nuclei if a collision were to eject portions of a neutron star, which then rapidly expands freed from confinement. That sequence could also begin earlier in galactic time than woulds-process nucleosynthesis; so each scenario fits the earlier growth ofr-process abundances in the galaxy. Each of these scenarios is the subject of active theoretical research. Observational evidence of the earlyr-process enrichment of interstellar gas and of subsequent newly formed stars, as applied to the abundance evolution of the galaxy of stars, was first laid out byJames W. Truranin 1981.[13]He and subsequent astronomers showed that the pattern of heavy-element abundances in the earliest metal-poor stars matched that of the shape of the solarr-process curve, as if thes-process component were missing. This was consistent with the hypothesis that thes-process had not yet begun to enrich interstellar gas when these young stars missing thes-process abundances were born from that gas, for it requires about 100 million years of galactic history for thes-process to get started whereas ther-process can begin after two million years. Theses-process–poor,r-process–rich stellar compositions must have been born earlier than anys-process, showing that ther-process emerges from quickly evolving massive stars that become supernovae and leave neutron-star remnants that can merge with another neutron star. The primary nature of the earlyr-process thereby derives from observed abundance spectra in old stars[4]that had been born early, when the galactic metallicity was still small, but that nonetheless contain their complement ofr-process nuclei.

Periodic tableshowing the cosmogenic origin of each element. The elements heavier than iron with origins in supernovae are typically those produced by ther-process, which is powered by supernova neutron bursts

Either interpretation, though generally supported by supernova experts, has yet to achieve a totally satisfactory calculation ofr-process abundances because the overall problem is numerically formidable. However, existing results are supportive; in 2017, new data about ther-process was discovered when theLIGOandVirgogravitational-wave observatories discovered a merger of two neutron stars ejectingr-process matter.[14]SeeAstrophysical sitesbelow.

Noteworthy is that ther-process is responsible for our natural cohort of radioactive elements, such as uranium and thorium, as well as the most neutron-rich isotopes of each heavy element.

Nuclear physics

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There are three natural candidate sites forr-process nucleosynthesis where the required conditions are thought to exist: low-masssupernovae,Type II supernovae,andneutron star mergers.[15]

Immediately after the severe compression of electrons in a Type II supernova,beta-minus decayis blocked. This is because the high electron density fills all available free electron states up to aFermi energywhich is greater than the energy of nuclear beta decay. However, nuclearcapture of those free electronsstill occurs, and causes increasingneutronizationof matter. This results in an extremely high density of free neutrons which cannot decay, on the order of 1024neutrons per cm3,[1]and hightemperatures.As this re-expands and cools,neutron captureby still-existing heavy nuclei occurs much faster thanbeta-minus decay.As a consequence, ther-process runs up along theneutron drip lineand 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-capturecross sectionin nuclei with closedneutron shells,the inhibiting process ofphotodisintegration,and the degree of nuclear stability in the heavy-isotope region. Neutron captures inr-process nucleosynthesis leads to the formation of neutron-rich,weakly boundnuclei withneutron separation energiesas low as 2 MeV.[16][1]At this stage, closed neutron shells atN= 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,[1]resulting in a slowdown orfreeze-outof the reaction.[17]

Decreasing nuclear stability terminates ther-process when its heaviest nuclei become unstable to spontaneous fission, when the total number of nucleons approaches 270. Thefission barriermay 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 unstableradioactivenuclei undergo a rapid succession of beta decays until they reach more stable, neutron-rich nuclei.[20]While thes-processcreates an abundance of stable nuclei having closed neutron shells, ther-process, in neutron-rich predecessor nuclei, creates an abundance of radioactive nuclei about 10amubelow thes-process peaks.[21]These abundance peaks correspond to stableisobarsproduced from successive beta decays of waiting point nuclei havingN= 50, 82, and 126—which are about 10 protons removed from theline of beta stability.[22]

Ther-process also occurs in thermonuclear weapons, and was responsible for the initial discovery of neutron-rich almost stable isotopes ofactinideslikeplutonium-244and the new elementseinsteiniumandfermium(atomic numbers 99 and 100) in the 1950s. It has been suggested that multiple nuclear explosions would make it possible to reach theisland of stability,as the affected nuclides (starting with uranium-238 as seed nuclei) would not have time to beta decay all the way to the quicklyspontaneously fissioningnuclides at the line of beta stability before absorbing more neutrons in the next explosion, thus providing a chance to reach neutron-richsuperheavynuclides likecopernicium-291 and -293 which may have half-lives of centuries or millennia.[23]

Astrophysical sites

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The most probable candidate site for ther-process has long been suggested to be core-collapsesupernovae(spectral typesIb,IcandII), which may provide the necessary physical conditions for ther-process. However, the very low abundance ofr-processnucleiin the interstellar gas limits the amount each can have ejected. It requires either that only a small fraction of supernovae ejectr-process nuclei to theinterstellar medium,or that each supernova ejects only a very small amount ofr-process material. The ejected material must be relatively neutron-rich, a condition which has been difficult to achieve in models,[2]so that astrophysicists remain uneasy about their adequacy for successfulr-process yields.

In 2017, new astronomical data about ther-process was discovered in data from the merger of twoneutron stars.Using the gravitational wave data captured inGW170817to identify the location of the merger, several teams[24][25][26]observed and studied optical data of the merger, finding spectroscopic evidence ofr-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 radioactiver-process matter of lower-mass-range heavy nuclei (A< 140such asstrontium)[27]and cooler red masses of higher mass-numberr-process nuclei (A> 140) rich inactinides(such asuranium,thorium,andcalifornium). 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 byr-process nuclei near their waiting points. Two distinct mass regions (A< 140andA> 140) for ther-process yields have been known since the first time dependent calculations of ther-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.[3]

These results offer a new possibility for clarifying six decades of uncertainty over the site of origin ofr-process nuclei. Confirming relevance to ther-process is that it is radiogenic power from radioactive decay ofr-process nuclei that maintains the visibility of these spun offr-process fragments. Otherwise they would dim quickly. Such alternative sites were first seriously proposed in 1974[28]as decompressingneutron starmatter. It was proposed such matter is ejected fromneutron starsmerging withblack holesin compact binaries. In 1989[29](and 1999[30]) this scenario was extended to binaryneutron starmergers (abinary star systemof two neutron stars that collide). After preliminary identification of these sites,[31]the scenario was confirmed inGW170817.Current astrophysical models suggest that a single neutron star merger event may have generated between 3 and 13Earth massesof gold.[32]

See also

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Notes

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  1. ^neutrons 1,674,927,471,000,000,000,000,000/cc vs 1 atom/ccinterstellar space
  2. ^Neutron number50, 82 and 126
  3. ^Abundance peaks for ther- ands-processes are atA= 80, 130, 196 andA= 90, 138, 208, respectively.

References

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  1. ^abcdef Burbidge, E. M.; Burbidge, G. R.; Fowler, W. A.; Hoyle, F. (1957)."Synthesis of the Elements in Stars".Reviews of Modern Physics.29(4): 547–650.Bibcode:1957RvMP...29..547B.doi:10.1103/RevModPhys.29.547.
  2. ^ab Thielemann, F.-K.;et al. (2011). "What are the astrophysical sites for ther-process and the production of heavy elements? ".Progress in Particle and Nuclear Physics.66(2): 346–353.Bibcode:2011PrPNP..66..346T.doi:10.1016/j.ppnp.2011.01.032.S2CID119412716.
  3. ^ab 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".Nature.551(7678): 80–84.arXiv:1710.05463.Bibcode:2017Natur.551...80K.doi:10.1038/nature24453.PMID29094687.
  4. ^ab Frebel, A.; Beers, T. C. (2018)."The formation of the heaviest elements".Physics Today.71(1): 30–37.arXiv:1801.01190.Bibcode:2018PhT....71a..30F.doi:10.1063/pt.3.3815.Nuclear physicists are still working to model ther-process, and astrophysicists need to estimate the frequency of neutron-star mergers to assess whetherr-process heavy-element production solely or at least significantly takes place in the merger environment.
  5. ^Cowan, John J.; Thielemann, Friedrich-Karl Thielemann (2004)."R-Process Nucleosynthesis in Supernovae"(PDF).Physics Today.57(10): 47–54.Bibcode:2004PhT....57j..47C.doi:10.1063/1.1825268.
  6. ^ab Hoyle, F. (1946)."The Synthesis of the Elements from Hydrogen".Monthly Notices of the Royal Astronomical Society.106(5): 343–383.Bibcode:1946MNRAS.106..343H.doi:10.1093/mnras/106.5.343.
  7. ^Suess, H. E.; Urey, H. C. (1956). "Abundances of the Elements".Reviews of Modern Physics.28(1): 53–74.Bibcode:1956RvMP...28...53S.doi:10.1103/RevModPhys.28.53.
  8. ^Woosley, Stan;Trimble, Virginia;Thielemann, Friedrich-Karl (2019). "The origin of the elements".Physics Today.72(2): 36–37.Bibcode:2019PhT....72b..36W.doi:10.1063/PT.3.4134.S2CID186549912.
  9. ^ Cameron, A. G. W. (1957)."Nuclear reactions in stars and nucleogenesis".Publications of the Astronomical Society of the Pacific.69(408): 201.Bibcode:1957PASP...69..201C.doi:10.1086/127051.
  10. ^ab Seeger, P. A.; Fowler, W. A.; Clayton, D. D. (1965)."Nucleosynthesis of heavy elements by neutron capture".Astrophysical Journal Supplement.11:121–66.Bibcode:1965ApJS...11..121S.doi:10.1086/190111.
  11. ^SeeSeeger, Fowler & Clayton 1965.Figure 16 shows the short-flux calculation and its comparison with naturalr-process abundances whereas Figure 18 shows the calculated abundances for long neutron fluxes.
  12. ^See Table 4 inSeeger, Fowler & Clayton 1965.
  13. ^ Truran, J. W. (1981). "A new interpretation of the heavy-element abundances in metal-deficient stars".Astronomy and Astrophysics.97(2): 391–93.Bibcode:1981A&A....97..391T.
  14. ^ Abbott, B. P.; et al. (LIGO Scientific Collaboration and Virgo Collaboration) (2017)."GW170817: Observation of Gravitational Waves from a Binary Neutron Star Inspiral".Physical Review Letters.119(16): 161101.arXiv:1710.05832.Bibcode:2017PhRvL.119p1101A.doi:10.1103/PhysRevLett.119.161101.PMID29099225.
  15. ^Bartlett, A.; Görres, J.; Mathews, G.J.; Otsuki, K.; Wiescher, W. (2006)."Two-neutron capture reactions and therprocess "(PDF).Physical Review C.74(1): 015082.Bibcode:2006PhRvC..74a5802B.doi:10.1103/PhysRevC.74.015802.Archived fromthe original(PDF)on 2020-08-06.Retrieved2019-06-17.
  16. ^Thoennessen, M. (2004)."Reaching the limits of nuclear stability"(PDF).Reports on Progress in Physics.67(7): 1187–1232.Bibcode:2004RPPh...67.1187T.doi:10.1088/0034-4885/67/7/R04.S2CID250790169.
  17. ^abEichler, M.A. (2016).Nucleosynthesis in explosive environments: neutron star mergers and core-collapse supernovae(PDF)(Doctoral thesis). University of Basel.
  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.Bibcode:2015PhRvC..92c1303W.doi:10.1103/PhysRevC.92.031303.S2CID59020556.
  19. ^ Boleu, R.; Nilsson, S. G.; Sheline, R. K. (1972)."On the termination of ther-process and the synthesis of superheavy elements ".Physics Letters B.40(5): 517–521.Bibcode:1972PhLB...40..517B.doi:10.1016/0370-2693(72)90470-4.
  20. ^ Clayton, D. D.(1968),Principles of Stellar Evolution and Nucleosynthesis,Mc-Graw-Hill, pp.577–91,ISBN978-0226109534,provides a clear technical introduction to these features. A more technical description can be found inSeeger, Fowler & Clayton 1965.
  21. ^Figure 10 ofSeeger, Fowler & Clayton 1965shows this path of captures reaching magic neutron numbers 82 and 126 at smaller values of nuclear charge Z than it does along the stability path.
  22. ^Surman, R.; Mumpower, M.; Sinclair, R.; Jones, K. L.; Hix, W. R.; McLaughlin, G. C. (2014)."Sensitivity studies for the weak r process: neutron capture rates".AIP Advances.4(41008): 041008.Bibcode:2014AIPA....4d1008S.doi:10.1063/1.4867191.
  23. ^ Zagrebaev, V.; Karpov, A.; Greiner, W. (2013)."Future of superheavy element research: Which nuclei could be synthesized within the next few years?".Journal of Physics: Conference Series.420(1): 012001.arXiv:1207.5700.Bibcode:2013JPhCS.420a2001Z.doi:10.1088/1742-6596/420/1/012001.
  24. ^ Arcavi, I.; et al. (2017)."Optical emission from a kilonova following a gravitational-wave-detected neutron-star merger".Nature.551(7678): 64–66.arXiv:1710.05843.Bibcode:2017Natur.551...64A.doi:10.1038/nature24291.
  25. ^Pian, E.; et al. (2017)."Spectroscopic identification ofr-process nucleosynthesis in a double neutron-star merger ".Nature.551(7678): 67–70.arXiv:1710.05858.Bibcode:2017Natur.551...67P.doi:10.1038/nature24298.PMID29094694.
  26. ^ Smartt, S. J.; et al. (2017)."A kilonova as the electromagnetic counterpart to a gravitational-wave source".Nature.551(7678): 75–79.arXiv:1710.05841.Bibcode:2017Natur.551...75S.doi:10.1038/nature24303.PMID29094693.
  27. ^Watson, Darach; Hansen, Camilla J.; Selsing, Jonatan; Koch, Andreas; Malesani, Daniele B.; Andersen, Anja C.; Fynbo, Johan P. U.;Arcones, Almudena;Bauswein, Andreas; Covino, Stefano; Grado, Aniello (2019). "Identification of strontium in the merger of two neutron stars".Nature.574(7779): 497–500.arXiv:1910.10510.Bibcode:2019Natur.574..497W.doi:10.1038/s41586-019-1676-3.ISSN0028-0836.PMID31645733.S2CID204837882.
  28. ^ Lattimer, J. M.; Schramm, D. N. (1974)."Black Hole–Neutron Star Collisions".Astrophysical Journal Letters.192(2): L145–147.Bibcode:1974ApJ...192L.145L.doi:10.1086/181612.
  29. ^ Eichler, D.; Livio, M.; Piran, T.; Schramm, D. N. (1989)."Nucleosynthesis, neutrino bursts and gamma-rays from coalescing neutron stars".Nature.340(6229): 126–128.Bibcode:1989Natur.340..126E.doi:10.1038/340126a0.
  30. ^ Freiburghaus, C.; Rosswog, S.; Thielemann, F.-K (1999)."r-process in Neutron Star Mergers ".Astrophysical Journal Letters.525(2): L121–L124.Bibcode:1999ApJ...525L.121F.doi:10.1086/312343.PMID10525469.
  31. ^ Tanvir, N.; et al. (2013)."A 'kilonova' associated with the short-duration gamma-ray burst GRB 130603B".Nature.500(7464): 547–9.arXiv:1306.4971.Bibcode:2013Natur.500..547T.doi:10.1038/nature12505.PMID23912055.
  32. ^"Neutron star mergers may create much of the universe's gold".Sid Perkins.Science AAAS. 20 March 2018.Retrieved24 March2018.
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