s-process

The slow neutron-capture process , or s-process is a series of reactions in nuclear astrophysics that occur in stars, particularly AGB stars. The s-process is responsible for the creation (nucleosynthesis) of approximately half the atomic nuclei heavier than iron.

In the s-process, a seed nucleus undergoes neutron capture to form an isotope with one higher atomic mass. If the new isotope is stable, a series of increases in mass can occur, but if it is unstable, then beta decay will occur, producing an element of the next highest atomic number. The process is slow (hence the name) in the sense that there is sufficient time for this radioactive decay to occur before another neutron is captured. A series of these reactions produces stable isotopes by moving along the valley of beta-decay stable isobars in the table of nuclides.

A range of elements and isotopes can be produced by the s-process, because of the intervention of alpha decay steps along the reaction chain. The relative abundances of elements and isotopes produced depends on the source of the neutrons and how their flux changes over time. Each branch of the s-process reaction chain eventually terminates at a cycle involving lead, bismuth, and polonium.

The s-process contrasts with the r-process, in which successive neutron captures are rapid: they happen more quickly than the beta decay can occur. The r-process dominates in environments with higher fluxes of free neutrons; it produces heavier elements and more neutron-rich isotopes than the s-process. Together the two processes account for most of the relative abundance of chemical elements heavier than iron.

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, these 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 how 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, whose longest half-life is 4.2 million years, had been discovered in s-, M-, and N-type stars in 1952[2][3] by Paul W. Merrill.[4][5] 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, probably unconnected with the nuclear fusion in the deep interior of the star that provides its power.

Nucleosynthesis periodic table
Periodic table showing the cosmogenic origin of each element. The elements heavier than iron with origins in dying low-mass stars are typically those produced by the s-process, which is characterized by slow neutron diffusion and capture over long periods in such stars

A calculable model for creating the heavy isotopes from iron seed nuclei in a time-dependent manner was not provided until 1961.[6] 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 neutron flux (number of neutrons per unit area) was appropriate. It also showed that no one single value for neutron flux could account for the observed s-process abundances, but that a wide range is required. The numbers of iron seed nuclei that were exposed to a given flux must decrease as the flux 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, as B2FH had sketched, but rather has a ledge-precipice structure. A series of papers[7][8][9][10][11][12] in the 1970s by Donald D. Clayton utilizing an exponentially declining neutron flux as a function of the number of iron seed 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[13] and by Karlsruhe Nuclear Physics Center in 1982[14] 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, seeded by iron nuclei left by a supernova during a previous generation of 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, passing decades between neutron captures. 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. The quantitative yield is also proportional to 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
6
C
 
4
2
He
 
→  16
8
O
 

n
22
10
Ne
 
4
2
He
 
→  25
12
Mg
 

n
S-process-elem-Ag-to-Sb
The s-process acting in the range from Ag to Sb.

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 sites of the main component are low-mass asymptotic giant branch stars.[15] The main component relies on the 13C neutron source above.[16] The weak component of the s-process, on the other hand, synthesizes s-process isotopes of elements from iron group seed nuclei to 58Fe on up to Sr and Y, and takes place at the end of helium- and carbon-burning in massive stars. It employs primarily the 22Ne neutron source. These stars will become supernovae at their demise and spew those s-process isotopes into interstellar gas.

The s-process is sometimes approximated over a small mass region using the so-called "local approximation", by which the ratio of abundances is inversely proportional to the ratio of neutron-capture cross-sections for nearby isotopes on the s-process path. This approximation is – as the name indicates – only valid locally, meaning for isotopes of nearby mass numbers, but it is invalid at magic numbers where the ledge-precipice structure dominates.

S-R-processes-atomic-mass-201-to-210
Chart representing the final part of the s-process. Red horizontal lines with a circle in their right ends represent neutron captures; blue arrows pointing up-left represent beta decays; green arrows pointing down-left represent alpha decays; cyan arrows pointing down-right represent electron captures.

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
83
Bi
 

n
 
→  210
83
Bi
 

γ
210
83
Bi
 
    →  210
84
Po
 

e
 

ν
e
210
84
Po
 
    →  206
82
Pb
 
4
2
He

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

206
82
Pb
 

n
 
→  209
82
Pb
209
82
Pb
 
    →  209
83
Bi
 
 
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:

   
n
 
→  4
2
He
 

e
 

ν
e
 

γ

The process thus terminates in bismuth, the heaviest "stable" element, and polonium, the first non-primordial element after bismuth. 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 short-lived polonium decays with a half-life of 138 days to stable lead.

The s-process measured in stardust

Stardust is one component of cosmic dust. Stardust is individual solid grains that condensed during mass loss from various long-dead stars. Stardust existed throughout interstellar gas before the birth of the Solar System and was trapped in meteorites when they assembled from interstellar matter contained in the planetary accretion disk in early Solar System. Today they are found in meteorites, where they have been preserved. Meteoriticists habitually refer to them as presolar grains. The s-process enriched grains are mostly silicon carbide (SiC). The origin of these grains is demonstrated by laboratory measurements of extremely unusual isotopic abundance ratios within the grain. First experimental detection of s-process xenon isotopes was made in 1978,[17] confirming earlier predictions that s-process isotopes would be enriched, nearly pure, in stardust from red giant stars.[18] These discoveries launched new insight into astrophysics and into the origin of meteorites in the Solar System.[19] Silicon carbide (SiC) grains condense in the atmospheres of AGB stars and thus trap isotopic abundance ratios as they existed in that star. Because the AGB stars are the main site of the s-process in the galaxy, the heavy elements in the SiC grains contain almost pure s-process isotopes in elements heavier than iron. This fact has been demonstrated repeatedly by sputtering-ion mass spectrometer studies of these stardust presolar grains.[19] Several surprising results have shown that within them 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 AGB-star atmospheres changed with time or from star to star, presumably with the strength of neutron flux in that star or perhaps the temperature. This is a frontier of s-process studies today.

References

  1. ^ 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. ^ Hammond, C. R. (2004). "The Elements". Handbook of Chemistry and Physics (81st ed.). CRC Press. ISBN 978-0-8493-0485-9.
  3. ^ Moore, C. E. (1951). "Technetium in the Sun". Science. 114 (2951): 59–61. Bibcode:1951Sci...114...59M. doi:10.1126/science.114.2951.59. PMID 17782983.
  4. ^ Merrill, P. W. (1952). "Technetium in the stars". Science. 115 (2992): 484.
  5. ^ George Sivulka (8 March 2017). "An Introduction to the Evidence for Stellar Nucleosynthesis". Stanford University. Retrieved 3 May 2018.
  6. ^ Clayton, D. D.; Fowler, W. A.; Hull, T. E.; Zimmerman, B. A. (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.
  7. ^ Clayton, D. D.; Rassbach, M. E. (1967). "Termination of the s-process". The Astrophysical Journal. 148: 69. Bibcode:1967ApJ...148...69C. doi:10.1086/149128.
  8. ^ Clayton, D. D. (1968). "Distribution of neutron-source strengths for the s-process". In Arnett, W. D.; Hansen, C. J.; Truran, J. W.; Cameron, A. G. W. (eds.). Nucleosynthesis. Gordon and Breach. pp. 225–240.
  9. ^ Peters, J. G.; Fowler, W. A.; Clayton, D. D. (1972). "Weak s-process Irradiations". The Astrophysical Journal. 173: 637. Bibcode:1972ApJ...173..637P. doi:10.1086/151450.
  10. ^ Clayton, D. D.; Newman, M. J. (1974). "s-process Studies: Exact Solution to a Chain Having Two Distinct Cross-Section Values". The Astrophysical Journal. 192: 501. Bibcode:1974ApJ...192..501C. doi:10.1086/153082.
  11. ^ Clayton, D. D.; Ward, R. A. (1974). "s-process Studies: Exact Evaluation of an Exponential Distribution of Exposures". The Astrophysical Journal. 193: 397. Bibcode:1974ApJ...193..397C. doi:10.1086/153175.
  12. ^ Ward, R. A.; Newman, M. J.; Clayton, D. D. (1976). "s-process Studies: Branching and the Time Scale". The Astrophysical Journal Supplement Series. 31: 33. Bibcode:1976ApJS...31...33W. doi:10.1086/190373.
  13. ^ Macklin, R. L.; Gibbons, J. H. (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.
  14. ^ Kaeppeler, F.; Beer, H.; Wisshak, K.; Clayton, D. D.; Macklin, R. L.; Ward, R. A. (1982). "s-process studies in the light of new experimental cross sections". The Astrophysical Journal. 257: 821–846. Bibcode:1982ApJ...257..821K. doi:10.1086/160033.
  15. ^ Boothroyd, A. I. (2006). "Heavy elements in stars". Science. 314 (5806): 1690–1691. doi:10.1126/science.1136842. PMID 17170281.
  16. ^ Busso, M.; Gallino, R.; Wasserburg, G. J. (1999). "Nucleosynthesis in Asymptotic Giant Branch Stars: Relevance for Galactic Enrichment and Solar System Formation" (PDF). Annual Review of Astronomy and Astrophysics. 37 (1): 239–309. Bibcode:1999ARA&A..37..239B. doi:10.1146/annurev.astro.37.1.239.
  17. ^ Srinivasan, B.; Anders, E. (1978). "Noble Gases in the Murchison Meteorite: Possible Relics of s-process Nucleosynthesis". Science. 201 (4350): 51–56. Bibcode:1978Sci...201...51S. doi:10.1126/science.201.4350.51. PMID 17777755.
  18. ^ Clayton, D. D.; Ward, R. A. (1978). "s-process studies: Xenon and krypton isotopic abundances". The Astrophysical Journal. 224: 1000. Bibcode:1978ApJ...224.1000C. doi:10.1086/156449.
  19. ^ a b Clayton, D. D.; Nittler, L. R. (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.
1 Aurigae

1 Aurigae is the original name for a star now in the constellation Perseus. It was the first entry in John Flamsteed's catalogue of stars in Auriga. When Eugène Joseph Delporte drew up simplified boundaries for the constellations on behalf of the International Astronomical Union in 1930, 1 Aurigae ended up over the border in Perseus. To avoid confusion, the star may instead be referred to by its Harvard Revised catalogue number, HR 1533.

Based upon its parallax measurement of 6.48 mas, this star is located approximately 520 light years from Earth. It is visible to the naked eye as a faint, orange-hued star with an apparent visual magnitude of 4.89. 1 Aurigae is moving closer to the Earth with a heliocentric radial velocity of −25 km/s.This is a possible binary star system, based upon the status of the visible component as a mild barium star. The primary is an aging giant star with a stellar classification of K3.5 III Ba0.2:. It is 3.9 billion years old with 1.49 times the mass of the Sun and around 44 times the Sun's radius. This star is radiating 561 times the luminosity of the Sun from its enlarged photosphere at an effective temperature of 4,102 K. The suspected companion star should be a white dwarf that previously transferred s-process elements to the visible member.

24 Aquilae

24 Aquilae (abbreviated 24 Aql) is a star in the equatorial constellation of Aquila. 24 Aquilae is its Flamsteed designation. It is at a distance of around 410 light-years (130 parsecs) from Earth and has an apparent visual magnitude of 6.4. According to the Bortle Dark-Sky Scale, this star is just visible to the naked eye in dark rural skies.

This is a so-called mild barium star, as identified by the mild presence of an absorption line of singly-ionized barium atoms at a wavelength of 455.4 nm. Such stars display an atmospheric overabundance of carbon and the heavy elements produced by the s-process, which was most likely transferred into the atmosphere by a wide binary stellar companion. However, in the case of 24 Aquilae, the abundances of heavy elements are near normal.At an estimated age of a half billion years, 24 Aquilae is a giant star with a stellar classification of K0 IIIa. It has more than double the mass of the Sun and shines with 63 times the Sun's luminosity. It is radiating this energy into space from its outer atmosphere at an effective temperature of 4,810 K. This heat is what gives it the cool orange hue characteristic of a K-type star.

2 Aurigae

2 Aurigae is a possible binary star system in the northern constellation of Auriga. This object is visible to the naked eye as a faint, orange-hued star with an apparent visual magnitude of +4.79. It forms an attractive four-star asterism when viewed in a low power eyepiece, together with the nearby HIP 22647 and another very loose visual pair, HIP 22776 and HIP 22744, all above magnitude 8. 2 Aurigae is moving closer to the Earth with a heliocentric radial velocity of −17 km/s.The visible component is an aging giant star with a stellar classification of K3- III Ba0.4. The suffix notation indicates this is a mild barium star, which means the stellar atmosphere is enriched with s-process elements. It is either a member of a close binary system and has previously acquired these elements from a (now) white dwarf companion or else it is on the asymptotic giant branch and is generating the elements itself. 2 Aurigae is 1.80 billion years old with 2.86 times the mass of the Sun and has expanded to 48 times the Sun's radius. It is radiating 599 times the Sun's luminosity from its enlarged photosphere at an effective temperature of 4,115 K.

58 Leonis

58 Leonis is a possible binary star system in the southern part of the constellation of Leo, near the border with Sextans. It shines with an apparent magnitude of 4.85, making it bright enough to be seen with the naked eye. An annual parallax shift of 9.05±0.20 mas yields a distance estimate of 360 light years. It is moving further from the Sun with a heliocentric radial velocity of +6 km/s.This orange hued star is an evolved K-type giant with a stellar classification of K0.5 III Fe-0.5, indicating a mild underabundance of iron in its spectrum. It was identified as a barium star by P. M. Williams (1971). These are theorized to be stars that show an enrichment of s-process elements by mass transfer from a now-white dwarf companion when it passed through the asymptotic giant branch stage. MacConnell et al. (1972) classified 58 Leonis as a marginal barium star. De Castro et al. (2016) consider this to be only a probable barium star, because of the low degree of s-process enrichment, and they rejected it from their sample. Rather than having an evolved companion, it may instead have formed from a cloud that was mildly enriched with s-process elements.

8 Andromedae

8 Andromedae, abbreviated 8 And, is a probable triple star system in the northern constellation of Andromeda. 8 Andromedae is the Flamsteed designation. It is visible to the naked eye with an apparent visual magnitude of 4.82. Based upon an annual parallax shift of 6.1 mas, it is located 540 light years from the Earth. It is moving closer with a heliocentric radial velocity of −8 km/s.The primary component is an aging red giant star with a stellar classification of M2.5 III Ba0.5. The suffix notation indicates this is a mild barium star, which means the stellar atmosphere is enriched with s-process elements. It is either a member of a close binary system and has previously acquired these elements from a (now) white dwarf companion or else it is on the asymptotic giant branch and is generating the elements itself. This is a periodic variable of unknown type, changing in brightness with an amplitude of 0.0161 magnitude at a frequency of 0.23354 d−1, or once every 4.3 days.The third component is a magnitude 13.0 star at an angular separation of 7.8″ along a position angle of 164°, as of 2015.

Barium star

Barium stars are spectral class G to K stars whose spectra indicate an overabundance of s-process elements by the presence of singly ionized barium, Ba II, at λ 455.4 nm. Barium stars also show enhanced spectral features of carbon, the bands of the molecules CH, CN and C2. The class was originally recognized and defined by William P. Bidelman and Philip Keenan. Initially, after their discovery, they were thought to be red giants, but the same chemical signature has been observed in main-sequence stars as well.

Observational studies of their radial velocity suggested that all barium stars are binary stars. Observations in the ultraviolet using International Ultraviolet Explorer detected white dwarfs in some barium star systems.Barium stars are believed to be the result of mass transfer in a binary star system. The mass transfer occurred when the now-observed giant star was on the main sequence. Its companion, the donor star, was a carbon star on the asymptotic giant branch (AGB), and had produced carbon and s-process elements in its interior. These nuclear fusion products were mixed by convection to its surface. Some of that matter "polluted" the surface layers of the main-sequence star as the donor star lost mass at the end of its AGB evolution, and it subsequently evolved to become a white dwarf. These systems are being observed at an indeterminate amount of time after the mass transfer event, when the donor star has long been a white dwarf. Depending on the initial properties of the binary system, the polluted star can be found at different evolutionary stages.During its evolution, the barium star will at times be larger and cooler than the limits of the spectral types G or K. When this happens, ordinarily such a star is spectral type M, but its s-process excesses may cause it to show its altered composition as another spectral peculiarity. While the star's surface temperature is in the M-type regime, the star may show molecular features of the s-process element zirconium, zirconium oxide (ZrO) bands. When this happens, the star will appear as an "extrinsic" S star.

Historically, barium stars posed a puzzle, because in standard stellar evolution theory G and K giants are not far enough along in their evolution to have synthesized carbon and s-process elements and mix them to their surfaces. The discovery of the stars' binary nature resolved the puzzle, putting the source of their spectral peculiarities into a companion star which should have produced such material. The mass transfer episode is believed to be quite brief on an astronomical timescale.

Prototypical barium stars include zeta Capricorni, HR 774, and HR 4474.

The CH stars are Population II stars with similar evolutionary state, spectral peculiarities, and orbital statistics, and are believed to be the older, metal-poor analogs of the barium stars.

Carbon star

A carbon star is typically an asymptotic giant branch star, a luminous red giant, whose atmosphere contains more carbon than oxygen. The two elements combine in the upper layers of the star, forming carbon monoxide, which consumes all the oxygen in the atmosphere, leaving carbon atoms free to form other carbon compounds, giving the star a "sooty" atmosphere and a strikingly ruby red appearance. There are also some dwarf and supergiant carbon stars, with the more common giant stars sometimes being called classical carbon stars to distinguish them.

In most stars (such as the Sun), the atmosphere is richer in oxygen than carbon. Ordinary stars not exhibiting the characteristics of carbon stars but cool enough to form carbon monoxide are therefore called oxygen-rich stars.

Carbon stars have quite distinctive spectral characteristics, and they were first recognized by their spectra by Angelo Secchi in the 1860s, a pioneering time in astronomical spectroscopy.

Dredge-up

A dredge-up is a period in the evolution of a star where a surface convection zone extends down to the layers where material has undergone nuclear fusion. As a result, the fusion products are mixed into the outer layers of the stellar atmosphere where they can appear in the spectrum of the star.

The first dredge-up occurs when a main-sequence star enters the red-giant branch. As a result of the convective mixing, the outer atmosphere will display the spectral signature of hydrogen fusion: the 12C/13C and C/N ratios are lowered, and the surface abundances of lithium and beryllium may be reduced.

The second dredge-up occurs in stars with 4–8 solar masses. When helium fusion comes to an end at the core, convection mixes the products of the CNO cycle. This second dredge-up results in an increase in the surface abundance of 4He and 14N, whereas the amount of 12C and 16O decreases.The third dredge-up occurs after a star enters the asymptotic giant branch and a flash occurs along a helium-burning shell. This dredge-up causes helium, carbon and the s-process products to be brought to the surface. The result is an increase in the abundance of carbon relative to oxygen, which can create a carbon star.The names of the dredge-ups are set by the evolutionary and structural state of the star in which each occurs, not by the sequence experienced by the star. As a result, lower-mass stars experience the first and third dredge-ups in their evolution but not the second.

Epsilon Aquilae

Epsilon Aquilae (ε Aql, ε Aquilae) is the Bayer designation for a binary star in the equatorial constellation of Aquila. It has an apparent visual magnitude of 4.02 and is visible to the naked eye. Based upon an annual parallax of 21.05 mas, Epsilon Aquilae lies at a distance of approximately 155 light-years (48 parsecs) from Earth. This is a spectroscopic binary system. The pair orbit each other over a period of 1,271 days (3.5 years) with an eccentricity of 0.27.It has the traditional name Deneb el Okab , from an Arabic term ذنب العقاب ðanab al-ʽuqāb "the tail of the eagle", and the Mandarin names Woo and Yuë , derived from and represent the state Wú (吳), an old state was located at the mouth of the Yangtze River, and Yuè (越), an old state in Zhejiang province (together with 19 Capricorni in Twelve States asterism). According to the R.H. Allen's works, it shares names with ζ Aquilae. Epsilon Aquilae is more precisely called Deneb el Okab Borealis, because is situated to the north of Zeta Aquilae, which can therefore be called Deneb el Okab Australis.

The primary component of this system is an evolved giant star with a stellar classification of K1 III. It has more than double the mass of the Sun and has expanded to ten times the Sun's radius. It shines with 54–fold the Sun's luminosity, which is being radiated from its outer envelope at an effective temperature of 4,760 K. At this heat, it glows with the orange-hue of a K-type star. This has been designated a barium star, meaning its atmosphere is extremely enriched with barium and other heavy elements. However, this is disputed, with astronomer Andrew McWilliam (1990) finding normal abundances from an s-process.

HD 122563

HD 122563 is an extremely metal-poor red giant star, and the brightest metal-poor star in the sky. Its low heavy element content was first recognized by spectroscopic analysis in 1963. For more than twenty years it was the most metal-poor star known, being more metal-poor than any known globular cluster, and it is the most accessible example of an extreme Population II or Halo star.

As the most extreme metal-poor star known, HD 122563's composition was crucial in constraining theories for galactic chemical evolution; in particular, its composition peculiarities provided signposts for understanding the accumulation of heavy elements by stellar nucleosynthesis in the Galaxy. For example, it has an excess of oxygen, [O/Fe] = +0.6, while the proportions of strontium, yttrium, zirconium, barium and the lanthanide elements suggest that the s-process has made no contribution to the material present in the star: in HD 122563, all these elements are products of the r-process instead. The implication is that the star formed at a time and place where there had not been enough time for any previous generation of stars to have produced s-process elements, though there was r-process material present.

Lead star

A lead star is a low-metallicity star with an overabundance of lead and bismuth as compared to other products of the S-process.

Metal

A metal (from Greek μέταλλον métallon, "mine, quarry, metal") is a material that, when freshly prepared, polished, or fractured, shows a lustrous appearance, and conducts electricity and heat relatively well. Metals are typically malleable (they can be hammered into thin sheets) or ductile (can be drawn into wires). A metal may be a chemical element such as iron, or an alloy such as stainless steel.

In physics, a metal is generally regarded as any substance capable of conducting electricity at a temperature of absolute zero. Many elements and compounds that are not normally classified as metals become metallic under high pressures. For example, the nonmetal iodine gradually becomes a metal at a pressure of between 40 and 170 thousand times atmospheric pressure. Equally, some materials regarded as metals can become nonmetals. Sodium, for example, becomes a nonmetal at pressure of just under two million times atmospheric pressure.

In chemistry, two elements that would otherwise qualify (in physics) as brittle metals—arsenic and antimony—are commonly instead recognised as metalloids, on account of their predominately non-metallic chemistry. Around 95 of the 118 elements in the periodic table are metals (or are likely to be such). The number is inexact as the boundaries between metals, nonmetals, and metalloids fluctuate slightly due to a lack of universally accepted definitions of the categories involved.

In astrophysics the term "metal" is cast more widely to refer to all chemical elements in a star that are heavier than the lightest two, hydrogen and helium, and not just traditional metals. A star fuses lighter atoms, mostly hydrogen and helium, into heavier atoms over its lifetime. Used in that sense, the metallicity of an astronomical object is the proportion of its matter made up of the heavier chemical elements.Metals comprise 25% of the Earth's crust and are present in many aspects of modern life. The strength and resilience of some metals has led to their frequent use in, for example, high-rise building and bridge construction, as well as most vehicles, many home appliances, tools, pipes, and railroad tracks. Precious metals were historically used as coinage, but in the modern era, coinage metals have extended to at least 23 of the chemical elements.The history of metals is thought to begin with the use of copper about 11,000 years ago. Gold, silver, iron (as meteoric iron), lead, and brass were likewise in use before the first known appearance of bronze in the 5th millennium BCE. Subsequent developments include the production of early forms of steel; the discovery of sodium—the first light metal—in 1809; the rise of modern alloy steels; and, since the end of World War II, the development of more sophisticated alloys.

Neutron capture

Neutron capture is a nuclear reaction in which an atomic nucleus and one or more neutrons collide and merge to form a heavier nucleus. Since neutrons have no electric charge, they can enter a nucleus more easily than positively charged protons, which are repelled electrostatically.Neutron capture plays an important role in the cosmic nucleosynthesis of heavy elements. In stars it can proceed in two ways: as a rapid (r-process) or a slow process (s-process). Nuclei of masses greater than 56 cannot be formed by thermonuclear reactions (i.e. by nuclear fusion), but can be formed by neutron capture.

Neutron capture on protons yields a line at 2.223 MeV predicted and commonly observed in solar flares.

Nu2 Sagittarii

Nu2 Sagittarii (ν2 Sagittarii) is a binary star system in the zodiac constellation of Sagittarius. It is faintly visible to the naked eye, having an apparent visual magnitude of +4.98. Based upon a small annual parallax shift of 11.91 mas as seen from Earth, this system is located around 270 light years from the Sun. Nu2 Sagittarii has a high peculiar velocity of 86.0+11.6−14.4 km/s and is most likely a runaway star system.The spectrum of the primary component displays a stellar classification of K1 Ib–II, indicating this is a K-type star with a mixed luminosity class of an evolved bright giant/supergiant star. It is a mild barium star, showing an enhanced abundance of s-process elements in its outer atmosphere. This material was most likely acquired during a previous mass transfer from its now white dwarf companion. The primary has an estimated 1.4 times the mass of the Sun and has expanded to 85 times the Sun's radius.

P-nuclei

p-nuclei (p stands for proton-rich) are certain proton-rich, naturally occurring isotopes of some elements between selenium and mercury inclusive which cannot be produced in either the s- or the r-process.

R-process

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". The other half are produced by the p-process and s-process. The r-process usually synthesizes all of the two most neutron-rich stable isotopes of each heavy element.

The heavy elements typically have six to ten stable isotopes. Chemical elements are defined by the number of protons in their atomic nucleus, e.g. all xenon atoms have 54 protons. But all elements also have neutrons in their atomic nucleus. Each isotope is characterized by the number of neutrons that it contains, e.g. xenon can have 70, 72, 74, 75, 76, 77, 78, 80, and 82 neutrons, and thus has 9 stable isotopes. The r-process contributes to the abundances of the heaviest four isotopes: 131Xe, 132Xe, 134Xe and 136Xe, and is solely responsible for the heaviest two of those. The s-process contributes to xenon's middle five isotopes: 128Xe, 129Xe, 130Xe, 131Xe, and 132Xe. The lightest two isotopes, 124Xe, and 126Xe, are produced by other processes.

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. The r-process abundances peak 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. This amounts to almost a gram of free neutrons in every cubic centimeter, an astonishing number requiring extreme locations. Traditionally this suggested the material ejected from the reexpanded core of a core-collapse supernova, as part of supernova nucleosynthesis, or decompression of neutron-star matter thrown off by a binary neutron star merger. The relative contributions of these sources to the astrophysical abundance of r-process elements is a matter of ongoing research.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.

The r-process contrasts with the s-process, the other predominant mechanism for the production of heavy elements, which is nucleosynthesis by means of slow captures of neutrons. The s-process primarily occurs within ordinary stars, particularly AGB stars, where the neutron flux is sufficient to cause neutron captures to recur every 10–100 years, much too slow for the r-process, which requires 100 captures per second. The s-process is secondary, 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. The r-process scenarios create their own seed nuclei, so they might proceed in massive stars that contain no heavy seed nuclei. Taken together, the r- and s-processes account for almost the entire abundance of chemical elements heavier than iron. The historical challenge has been to locate physical settings appropriate for their time scales.

Rp-process

The rp-process (rapid proton capture process) consists of consecutive proton captures onto seed nuclei to produce heavier elements. It is a nucleosynthesis process and, along with the s-process and the r-process, may be responsible for the generation of many of the heavy elements present in the universe. However, it is notably different from the other processes mentioned in that it occurs on the proton-rich side of stability as opposed to on the neutron-rich side of stability. The end point of the rp-process (the highest mass element it can create) is not yet well established, but recent research has indicated that in neutron stars it cannot progress beyond tellurium. The rp-process is inhibited by alpha decay, which puts an upper limit on the end point at 104Te, the lightest observed alpha decaying nuclide, and the proton drip line in light antimony isotopes. At this point, further proton captures result in prompt proton emission or alpha emission and thus the proton flux is consumed without yielding heavier elements; this end process is known as the tin-antimony-tellurium cycle.

S-type star

An S-type star (or just S star) is a cool giant with approximately equal quantities of carbon and oxygen in its atmosphere. The class was originally defined in 1922 by Paul Merrill for stars with unusual absorption lines and molecular bands now known to be due to s-process elements. The bands of zirconium monoxide (ZrO) are a defining feature of the S stars.

The carbon stars have more carbon than oxygen in their atmospheres. In most stars, such as class M giants, the atmosphere is richer in oxygen than carbon and they are referred to as oxygen-rich stars. S-type stars are intermediate between carbon stars and normal giants. They can be grouped into two classes: intrinsic S stars, which owe their spectra to convection of fusion products and s-process elements to the surface; and extrinsic S stars, which are formed through mass transfer in a binary system.

The intrinsic S-type stars are on the most luminous portion of the asymptotic giant branch, a stage of their lives lasting less than a million years. Many are long period variable stars. The extrinsic S stars are less luminous and longer-lived, often smaller-amplitude semiregular or irregular variables. S stars are relatively rare, with intrinsic S stars forming less than 10% of asymptotic giant branch stars of comparable luminosity, while extrinsic S stars form an even smaller proportion of all red giants.

Zeta Capricorni

Zeta Capricorni, Latinized from ζ Capricorni, is a fourth-magnitude star in the constellation Capricornus. ζ Capricorni is a binary star, with the primary component ζ Capricorni A being a yellow G-type giant with an apparent magnitude of +3.77. It is considered one of the prototypical examples of a Barium star, properties of which include overabundances of carbon molecules (such as C2) and s-process elements. Zeta Capricorni has an overabundance of the s-process element praseodymium.

Its companion, ζ Capricorni B is a hydrogen-rich white dwarf. It is about as massive as the Sun, and its temperature is 23,000 K.The ζ Capricorni binary system is approximately 390 light years from Earth, based on its parallax.

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