S-process

The S-process or slow-neutron-capture-process is a nucleosynthesis process that occurs at relatively low neutron density and intermediate temperature conditions in stars. Under these conditions the rate of neutron capture by atomic nuclei is slow relative to the rate of radioactive beta-minus decay. In the S-process, a stable isotope captures a neutron, but the radioactive isotope that results decays to its stable daughter before the next neutron is captured. This process produces stable isotopes by moving along the valley of beta-decay stable isobars in the chart of isotopes. The S-process produces approximately half of the isotopes of the elements heavier than iron, and therefore plays an important role in the galactic chemical evolution. The S-process differs from the more rapid R-process of neutron capture by its slow rate of neutron captures.

Contents 1 History 2 The S-process in stars 3 The S-process measured in stardust 4 References

History

The S-process was seen to be needed from the relative abundances of isotopes of heavy elements and from a newly published table of abundances by Hans Suess and Harold Urey in 1956. Among other things, this data showed abundance peaks for strontium, barium, and lead, which, according to quantum mechanics and the nuclear shell model, are particularly stable nuclei, much like the noble gases are chemically inert. This implied that some abundant nuclei must be created by slow neutron capture, and it was only a matter of determining what other nuclei could be accounted for by such a process. A table apportioning the heavy isotopes between S-process and R-process was published in the famous B²FH review paper in 1957.^[1] There it was also argued that the S-process occurs in red giant stars. In a particularly illustrative case, the element technetium, with a longest half-life of 4.2 million years, had been discovered in S-, M-, and N-type stars in 1952.^[2][3] Since these stars were thought to be billions of years old, the presence of technetium in their outer atmospheres was taken as evidence of its recent creation there, unconnected with events in the deep interior of the star in the region of active fusion, or events in the star's early history billions of years in the past.

A calculable model for creating the heavy isotopes from iron seed nuclei in a time-dependent manner was not provided until 1961.^[4] That work showed that the large overabundances of barium observed by astronomers in certain red-giant stars could be created from iron seed nuclei if the total fluence (number of neutrons per unit area) of neutrons was appropriate. It also showed that no one single fluence could account for the observed S-process abundances, but that a wide range of fluences is required. The numbers of iron seed nuclei that were exposed to a given fluence must decrease as the fluence becomes stronger. This work also showed that the curve of the product of neutron-capture cross section times abundance is not a smoothly falling curve, but rather has a ledge-precipice structure. A series of papers in the 1970s by D. Clayton based on the assumption of an exponentially declining neutron fluence as a function of the number of iron seed so exposed became the standard model of the S-process and remained so until the details of AGB-star nucleosynthesis became advanced enough that they became a standard model based on the stellar structure models. Important series of measurements of neutron-capture cross sections were reported from Oak Ridge National Lab in 1965^[5] and by Karlsruhe Nuclear Physics Center in 1982^[6] and subsequently. These placed the S-process on the firm quantitative basis that it enjoys today.

The S-process in stars

The S-process is believed to occur mostly in Asymptotic Giant Branch stars. In contrast to the R-process which is believed to occur over time scales of seconds in explosive environments, the S-process is believed to occur over time scales of thousands of years. The extent to which the s-process moves up the elements in the chart of isotopes to higher mass numbers is essentially determined by the degree to which the star in question is able to produce neutrons, and by the amount of iron in the star's initial abundance distribution. Iron is the "starting material" (or seed) for this neutron capture - beta-minus decay sequence of synthesizing new elements.

The main neutron source reactions are:

13 6C	+	4 2He	→	16 8O	+	n
22 10Ne	+	4 2He	→	25 12Mg	+	n

One distinguishes the main and the weak s-process component. The main component produces heavy elements beyond Sr and Y, and up to Pb in the lowest metallicity stars. The production site of the main component are low-mass Asymptotic Giant Branch stars.^[7] The weak component of the S-process, on the other hand, synthesizes S-process isotopes of elements from the iron group up to Sr and Y, and takes place at the end of He- and C-burning in massive stars. These stars will become supernovae at their demise and spew those s isotopes into interstellar space.

The S-process is often mathematically treated using the so-called local approximation, which gives a theoretical model of elemental abundances based on the assumption of constant neutron flux in a star, so that the ratio of abundances is inversely proportional to the ratio of neutron-capture cross-sections for different isotopes. This approximation is - as the name indicates - only valid locally, meaning for isotopes of similar mass number.

Because of the relatively low neutron fluxes expected to occur during the S-process (on the order of 10⁵ to 10¹¹ neutrons per cm² per second), this process does not have the ability to produce any of the heavy radioactive isotopes such as thorium or uranium. The cycle that terminates the S-process is:

209
Bi captures a neutron, producing 210
Bi, which decays to 210
Po by β^- decay. 210
Po in turn decays to 206
Pb by α decay:

209 83Bi	+	n	→	210 83Bi	+	γ
210 83Bi			→	210 84Po	+	e−	+	ν e
210 84Po			→	206 82Pb	+	4 2He

206
Pb then captures three neutrons, producing 209
Pb, which decays to 209
Bi by β^- decay, restarting the cycle:

206 82Pb	+	3 n	→	209 82Pb
209 82Pb			→	209 83Bi	+	e−	+	ν e

The net result of this cycle therefore is that 4 neutrons are converted into one alpha particle, two electrons, two anti-electron neutrinos and gamma radiation:

4 n

→

4 2He

+

2 e−

+

2 ν e

+

γ

The process thus terminates in bismuth, the heaviest "stable" element. (Bismuth is actually slightly radioactive, but with a half-life so long—a billion times the present age of the universe—that it is effectively stable over the lifetime of any existing star.)

The S-process measured in stardust

Stardust is one component of cosmic dust. Individual solid grains from various long-dead stars that existed before the solar system are found in meteorites, where they have been preserved. The origin of these grains is demonstrated by laboratory measurements of extremely unusual isotopic abundance ratios within the grain. The results give new insight into astrophysics.^[8] Silicon-carbide (SiC) grains condense in the atmospheres of AGB stars and thus trap the isotopes of that star. Because the AGB stars are the main site of the S-process in the galaxy, the heavy elements in the SiC grains are virtually pure S-process isotopes. This fact has been demonstrated repeatedly by sputtering-ion mass spectrometer studies of these presolar grains.^[8] Several surprising results have shown that the ratio of S-process and R-process abundances is somewhat different from that which was previously assumed. It has also been shown with trapped isotopes of krypton and xenon that the S-process abundances in the stellar atmospheres change with time or from star to star, presumably with the strength of neutron fluence or perhaps the temperature. This is a frontier of S-process studies today.

References

1. ^ E. M. Burbidge, G. R. Burbidge, W. A. Fowler, F. Hoyle (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. ^ Hammond, C. R. (2004). The Elements, in Handbook of Chemistry and Physics 81st edition. CRC press. ISBN 0-8493-0485-7. 
3. ^ Moore, CE (1951). "Technetium in the Sun.". Science 114 (2951): 59–61. Bibcode 1951Sci...114...59M. doi:10.1126/science.114.2951.59. PMID 17782983. 
4. ^ D. D. Clayton, W. A. Fowler, T. E. Hull, B. A. Zimmerman (1961). "Neutron capture chains in heavy element synthesis". Annals of Physics 12 (3): 331–408. Bibcode 1961AnPhy..12..331C. doi:10.1016/0003-4916(61)90067-7. 
5. ^ R. L. Macklin, J. H. Gibbons (1965). "Neutron Capture Data at Stellar Temperatures". Reviews of Modern Physics 37 (1): 166–176. Bibcode 1965RvMP...37..166M. doi:10.1103/RevModPhys.37.166. 
6. ^ F. Kaeppeler, H. Beer, K. Wisshak, D. D. Clayton, R. L. Macklin, R. A. Ward (1982). "S process studies in the light of new experimental cross sections". Astrophysical Journal 257: 821–846. Bibcode 1982ApJ...257..821K. doi:10.1086/160033. 
7. ^ A. I. Boothroyd (2006). "Heavy elements in stars". Science 314 (5806): 1690–1691. doi:10.1126/science.1136842. PMID 17170281. 
8. ^ ^{a b} D. D. Clayton, L. R. Nittler (2004). "Astrophysics with Presolar stardust". Annual Review of Astronomy and Astrophysics 42 (1): 39–78. Bibcode 2004ARA&A..42...39C. doi:10.1146/annurev.astro.42.053102.134022. 

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