p-cores

Certain proton-rich isotopes of the elements from selenium to mercury are referred to as p-nuclei (“p” for proton-rich ) , whose natural occurrence in the universe cannot be explained by the s- or r-process in the context of nucleosynthesis .

definition

Excerpt from the nuclide map with stable p-nuclei (yellow), r-nuclei (red) and s-nuclei (green)

According to the current theory of the formation of elements , of the atomic nuclei occurring in the universe, only the light nuclei (with atomic numbers up to that of iron ) can have been created by nuclear fusion . According to Burbidge, Burbidge, Fowler and Hoyle and AGW Cameron (both 1957), the presence of most naturally occurring nuclides beyond the element iron can be explained by two types of astrophysical neutron capture processes , the s-process and the r-process . However, some naturally occurring, proton-rich nuclides are not formed by these two processes . Therefore there has to be (at least) one other process by which these so-called p-cores arise.

Since the definition of the p-cores depends on our level of knowledge regarding the s and r processes, the original list of 35 p-cores may change over time, as indicated in the table below. Today it is assumed that the frequencies of ¹⁵² Gd and ¹⁶⁴ Er could at least contain strong contributions from the s-process . This also applies to those of the isotopes ¹¹³ In and ¹¹⁵ Sn, which could also be formed in small quantities in the r process.

The long-lived radionuclides ⁹² Nb, ⁹⁷ Tc, ⁹⁸ Tc and ¹⁴⁶ Sm do not belong to the classic p-nuclei because they do not occur on Earth. However, they are p-kernels by definition, since they are not generated in the s and r processes. By demonstrating their decay products in presolar grains it can be deduced that at least ⁹² Nb and ¹⁴⁶ Sm were present in traces in the primordial cloud of the solar system . This allows conclusions to be drawn about the length of time since the last production of p-nuclei before the formation of the solar system.

The p nuclei are very rare. Isotopes of an element that are p-nuclei are 10 to 1000 times rarer than the other isotopes of the element. The frequencies of the p-nuclei can only be determined by geochemical investigations and by analyzes of meteorites and presolar grains. They can not be identified in isolation in star spectra . Therefore the knowledge of the frequencies is limited to those in the solar system. It is therefore unclear whether the solar abundances of the p-nuclei are typical for the Milky Way .

List of p-kernels
nuclide	comment
⁷⁴ Se
⁷⁸ kr
⁸⁴ Sr
⁹² Nb	long-lived radionuclide ; not a classic p-kernel, but cannot be generated in s- or r-process
⁹² Mon
⁹⁴ months
⁹⁷ Tc	long-lived radionuclide ; not a classic p-kernel, but cannot be generated in s- or r-process
⁹⁸ Tc	long-lived radionuclide ; not a classic p-kernel, but cannot be generated in s- or r-process
⁹⁶ Ru
⁹⁸ Ru
¹⁰² Pd
¹⁰⁶ Cd
¹⁰⁸ Cd
¹¹³ in	Generated entirely or partially in the s-process ? Also contributions from the r-process ?
¹¹² Sn
¹¹⁴ Sn
¹¹⁵ Sn	Generated entirely or partially in the s-process ? Also contributions from the r-process ?
¹²⁰ te
¹²⁴ Xe
¹²⁶ Xe
¹³⁰ Ba
¹³² Ba
¹³⁸ La	Generated by ν process
¹³⁶ Ce
¹³⁸ Ce
¹⁴⁴ Sm
¹⁴⁶ Sm	long-lived radionuclide ; not a classic p-kernel, but cannot be generated in s- or r-process
¹⁵² Gd	generated entirely or partially in the s-process ?
¹⁵⁶ Dy
¹⁵⁸ Dy
¹⁶² he
¹⁶⁴ he	generated entirely or partially in the s-process ?
¹⁶⁸ Yb
¹⁷⁴ Hf
¹⁸⁰ days	partially generated in the ν process; also contribution from the s-process ?
¹⁸⁰ W
¹⁸⁴ Os
¹⁹⁰ pt
¹⁹⁶ ed

Origin of the p-nuclei

The mechanism by which the p-nuclei are generated is not yet fully understood. The previously favored γ-process in core collapse supernovae can not generate all p-nuclei to a sufficient extent in astrophysical model calculations . Therefore, further possibilities are being investigated, as explained below. It is also possible that there is not just a single responsible nucleosynthesis process, but that the various p-nuclei are generated by different processes in diverse astrophysical environments.

Principal Mechanisms

In principle, proton- rich nuclides can be generated in two ways: by sequences of proton captures (these are (p, γ) reactions ) on atomic nuclei with a lower number of protons or by sequences of photo - disintegration of neutron-rich atomic nuclei, which by means of (γ, n) -uclear reactions also have atomic nuclei generate a higher proportion of protons.

p-nuclei are difficult to obtain by capturing protons, because the higher the number of protons in the atomic nucleus, the higher the Coulomb wall , which each newly added proton has to overcome. The higher the Coulomb Wall, the more energy a proton needs to penetrate the nucleus and be captured there. The mean energy of the protons is determined by the temperature of the stellar plasma . If the temperature increases, however, the photodisintegrations are also faster and protons are knocked out of the atomic nuclei via (γ, p) -nuclear reactions than they can be deposited via (p, γ). The solution would be to have a large number of protons, so that the effective number of captures per second is large, even if the temperature is not increased too much. In extreme cases, this leads to the formation of very short-lived radionuclides, which only decay into stable nuclides after the process has been completed .

In the search for production possibilities for p-nuclei one therefore has to investigate suitable combinations of temperature and proton density of the stellar plasma. Further parameters are the time span in which the nuclear reactions can take place, as well as the number and type of nuclides that are assumed ( seed cores ).

Possible processes

The p-process

In the p-process one tries to produce p-nuclei by capturing a few protons on the corresponding nuclides from the s- and r-process, which must be in the stellar plasma from the start. As explained above, this process is not suitable for the generation of p-nuclei, although it was originally proposed for this. Sometimes the term p-process is used very generally for any process that generates p-kernels.

The γ process

p-nuclei can also be obtained by photo-disintegration of nuclides from the s- and r-process. At temperatures of 2 to 3 Giga kelvin and a short period of time (which requires an explosive process), the photo-disintegration takes place only very incompletely, just enough to generate the required small amounts of p-nuclei. Since the photodisintegration takes place via nuclear reactions of the types (γ, n), (γ, α) and (γ, p) triggered by high-energy photons ( gamma radiation ), this is called the γ process .

The ν process

Neutrino- induced nuclear reactions can directly produce certain nuclides, for example in nuclear collapse supernovae ⁷ Li, ¹¹ B, ¹⁹ F, ¹³⁸ La. This is called the neutrino or ν process and requires the existence of a neutrino source of sufficient intensity.

Rapid proton captures

In the p-process, protons are attached to stable or weakly radioactive atomic nuclei. However, if a high proton density is available, even very short-lived atomic nuclei will still capture one or more protons before they decay . As a result, the nucleosynthesis process quickly moves away from the stable nuclei to the proton-rich side of the nuclide map. This is called rapid proton capture .

A sequence of (p, γ) -nuclear reactions continues until either the β-decay of a nuclide is faster than a proton capture on it or it is no longer energetically favorable to add another proton (see the boundaries of the nuclide map defined by the binding energies ). In both cases one or more β-decays occur until a nucleus is generated which, under the given conditions, again captures protons faster than it decays. Then the reaction path continues.

Since, under sufficiently proton-rich conditions, both the proton capture and the decays are fast, one can get from the lightest nuclei up to ⁵⁶ Ni within a second . From there there are a number of waiting points in the response path , i. H. Nuclides that both have a long half-life and only slowly add an additional proton (i.e. their cross-section for (p, γ) reactions is also small). Examples of such waiting points are: ⁵⁶ Ni, ⁶⁰ Zn, ⁶⁴ Ge, ⁶⁸ Se. Depending on the position of the reaction path, other waiting points can also play a role. It is typical for waiting points that they have half-lives of minutes to days and thus drastically increase the time required to continue the reaction sequences. If the conditions required for rapid proton capture are only present for a short time (in explosive scenarios the time is typically in the range of a few seconds), the waiting points limit or slow down the continuation of the reactions to heavier nuclei (which, however, are still formed with a certain probability can be, even if to a small extent).

In order to generate p-nuclei, the process path must be continued to such an extent that nuclides with the mass numbers of the corresponding p-nuclei (but containing more protons) are formed. After the rapid proton capture has subsided, these are finally converted into the p nuclei via sequences of β decays.

The fast proton capture category includes the rp, pn, and νp processes, which are discussed below.

The rp process

The so-called rp process ( rp for rapid proton capture , dt. "Rapid proton capture ") is a pure, rapid proton capture process , as described above. With proton densities of more than 10 ²⁸ protons / cm³ and temperatures around 2 gigakelvin, the reaction path moves near the proton-rich limit of the nuclide map. The waiting points can be overcome with a correspondingly long period of 10 to 600 s. The waiting point nuclides are generated more frequently, while the formation of nuclei behind each waiting point is more and more suppressed.

A definite end point is reached in the area around ¹⁰⁷ Te, because there the reaction path runs into a region of nuclides, which preferentially convert through α-decay and thus repeatedly throw the reaction path back on itself in a loop. Thus, an rp process could only generate p-kernels with mass numbers A ≤ 107.

The pn process

The waiting points in fast proton capture can be bridged by (n, p) -nuclear reactions, which take place on waiting-point nuclei faster than proton capture and decays. This significantly reduces the time it takes to assemble heavy elements and enables efficient production in a few seconds. The neutrons required are released in other nuclear reactions.

The νp process

Another way to get the neutrons needed for (n, p) reactions is in proton-rich environments in the reaction that creates a positron and a neutron from an antineutrino and a proton. Since (anti-) neutrinos only interact very weakly with protons, a high proton density is required on which a correspondingly high flow of antineutrinos acts. ${\ displaystyle \ mathrm {{\ bar {\ nu}} _ {e} + p \ rightarrow e ^ {+} + n}}$

Possible places of synthesis

Core collapse supernovae

Massive stars end their lives as core collapse supernovae. An explosion shock wave runs from the center of the exploding star through the outer layers and blows them off. If this shock wave reaches the O / Ne shell of the star, the conditions for a γ process are met there for 1–2 seconds.

Although the majority of the p-nuclei can be generated in this way, some mass number ranges cause difficulties in model calculations of such supernovae. It has been known for a long time that the p-nuclei with mass numbers A <100 do not arise in the γ-process. Modern model calculations also show problems in the mass number range 150 ≤ A ≤ 165.

The p-core ¹³⁸ La does not arise in the γ process, but can be formed in the ν process. In such a supernova, a hot neutron star is formed inside , which emits neutrinos with high intensity . The neutrinos interact with the outer layers of the exploding star and cause nuclear reactions that may a. ^{138 Create} La. The natural ¹⁸⁰ Ta occurrences can also partly originate from this ν process.

It was proposed to supplement the γ-process taking place in the outer layers with another process that takes place in the deepest layers, closest to the neutron star, which are just about to be ejected. Due to the initially strong neutrino flux near the neutron star, these layers become extremely proton-rich (due to the reaction ν _e + n → e ^- + p). Although the antineutrino flow is weaker, neutrons can still be produced due to the high proton density and thus enable the νp process . Because of the very limited time span of the explosion and the high Coulomb wall of the heavier nuclei, the νp process could at most produce the lightest p-nuclei. How much of it it generates, however, depends very sensitively on many details of the simulations and on the as yet not fully understood explosion mechanism of the core collapse supernovae.

Thermonuclear supernovae

Thermonuclear supernovae are the explosions of white dwarfs triggered by the accretion of matter from the outer shell of a companion star on the surface of the white dwarf. The accreted material is rich in hydrogen (i.e. protons) and helium ( α-particles ) and becomes so hot that nuclear reactions start.

It is assumed that there are (at least) the two types of such explosions described below. In none of them are neutrinos released, which is why neither the ν nor the νp process can take place. The prerequisites for the rp process are also not met.

The details of the possible generation of p-nuclei in such supernovae depend sensitively on the composition of the material deposited by the companion star. Since this can fluctuate considerably from star to star, all statements and models on the formation of the p-nucleus in thermonuclear supernovae are subject to the corresponding uncertainties.

Type Ia supernovae

In the standard model of thermonuclear supernovae , the white dwarf explodes after it exceeds the Chandrasekhar mass by accretion , because this ignites explosive carbon burning under degenerate conditions . A nuclear burning front runs through the white dwarf from the inside out and tears it apart. In the outermost layers just below the surface of the white dwarf (which still contain 0.05 solar masses of material) the right conditions for a γ process then arise.

The p-nuclei are created in the same way as in the γ-process that takes place in nuclear collapse supernovae, and the same problems arise. In addition, ¹³⁸ La and ¹⁸⁰ Ta are not generated either. A variation of the seed core frequency by means of assumed increased frequencies of the nuclides formed in the s-process only scales the resulting frequencies of the p-nuclei, but cannot solve the problems of relative underproduction in the above-mentioned areas.

Sub-Chandrasekhar supernovae

In a subclass of Type Ia supernovae, called sub-Chandrasekhar supernovae , the white dwarf may explode before it reaches the Chandrasekhar limit because nuclear reactions occurring during accretion heat the white dwarf and trigger the ignition of the explosive carbon burner prematurely. This is favored if the accreted material is very rich in helium. At the bottom of the helium layer ignites helium burning explosive because of the degeneration and solves two shock fronts. The one running inwards reaches the center of the white dwarf and triggers the carbon explosion there. The one running outwards heats the outer layers of the white dwarf and repels them. In these outer layers, a γ process takes place again at temperatures of 2 to 3 gigakelvin. The presence of α-particles also enables nuclear reactions which release a large number of neutrons. These are, for example, the reactions ¹⁸ O (α, n) ²¹ Ne, ²² Ne (α, n) ²⁵ Mg and ²⁶ Mg (α, n) ²⁹ Si. This means that a pn process can also take place in that part of the outer layer whose temperature exceeds 3 gigakelvin .

The light p-nuclei, which are under-produced in the γ-process, can be generated so efficiently in the pn-process that they even achieve much higher frequencies than the remaining p-nuclei. In order to correct the relative frequencies, an increased s-process seed kernel frequency in the accreted material must be assumed, which then increases the yield of the γ-process.

Neutron stars in binary systems

A neutron star can also accumulate material from a companion star on its surface. If the accreted layer reaches a density of 10 ⁵ -10 ⁶ g / cm³ and a temperature of more than 0.2 gigakelvin, combined hydrogen and helium burning ignites under degenerate conditions. Similar to a sub-Chandrasekhar supernova, this leads to a thermonuclear explosion, which, however, cannot harm the neutron star. The nuclear reactions can thus take longer in the accreted layer than in an explosion, and an rp process results. This continues until either no more protons are present or until the density of the layer has become too low for the nuclear reactions due to the increase in temperature.

X-ray flashes observed in the Milky Way can be explained by an rp process on the surface of neutron stars. However, it remains unclear whether and how much matter can escape from the gravitational field of the neutron star. Only then would these objects be a possible source of p-nuclei. Even if this is the case, only the light p-nuclei (which are underproduced in core collapse supernovae) can arise because of the end point of the rp process.

Web links

Nuclear Astrophysics: Element Synthesis in the Universe .
The p-process . ( Memento from December 20, 2006 in the Internet Archive )

Individual evidence

^ ^A ^b E. M. Burbidge , GR Burbidge , WA Fowler , Fred Hoyle : Synthesis of the Elements in Stars . In: Reviews of Modern Physics . 29, No. 4, 1957, pp. 547-650. doi : 10.1103 / RevModPhys.29.547 .
^ ^A ^b A. GW Cameron : Nuclear Reactions in Stars and Nucleogenesis. In: Publications of the Astronomical Society of the Pacific. Vol. 69, 1957, pp. 201-222, bibcode : 1957PASP ... 69..201C .
↑ C. Arlandini, F. Käppeler, K. Wisshak, R. Gallino, M. Lugaro, M. Busso, O. Straniero: Neutron Capture in Low-Mass Asymptotic Giant Branch Stars: Cross Sections and Abundance Signatures. In: The Astrophysical Journal . Vol. 525, 1999, pp. 886-900. doi: 10.1086 / 307938 .
↑ Zs. Nemeth, F. Käppeler, C. Theis, T. Belgya, SW Yates: Nucleosynthesis in the Cd-In-Sn region. In: The Astrophysical Journal. Vol. 426, 1994, pp. 357-365. doi: 10.1086 / 174071 .
↑ N. Dauphas, T. Rauscher, B. Marty, L. Reisberg: Short-lived p-nuclides in the early solar system and implications on the nucleosynthetic role of X-ray binaries. In: Nuclear Physics. Vol. A719, 2003, pp. C287-C295, doi: 10.1016 / S0375-9474 (03) 00934-5 , arxiv : astro-ph / 0211452 .
↑ ^a ^b ^c ^d ^e ^f ^g ^h ⁱ ^j M. Arnould, S. Goriely: The p-process of stellar nucleosynthesis: astrophysics and nuclear physics status. In: Physics Reports. 384, 2003, pp. 1-84.
↑ ^a ^b ^c ^d T. Rauscher: Origin of p-Nuclei in Explosive Nucleosynthesis. In: Proceedings of Science. arxiv : 1012.2213v1 .
^ ^A ^b ^c S. E. Woosley, WM Howard: The p-process in supernovae. In: The Astrophysical Journal Supplement. Vol. 36, 1978, pp. 285-304. doi: 10.1086 / 190501 .
^ ^A ^b S. E. Woosley, DH Hartmann, RD Hoffman, WC Haxton: The ν-process. In: The Astrophysical Journal. Vol. 356, 1990, pp. 272-301. doi: 10.1086 / 168839 .
↑ ^a ^b ^c H. Schatz u. a .: rp-Process Nucleosynthesis at Extreme Temperature and Density Conditions. In: Physics Reports. Vol. 294, 1998, 167-263. doi: 10.1016 / S0370-1573 (97) 00048-3 .
↑ ^a ^b H. Schatz u. a .: End Point of the rp Process on Accreting Neutron Stars. In: Physical Review Letters. Vol. 86, 2001, pp. 3471-3474. doi: 10.1103 / PhysRevLett.86.3471 .
↑ ^a ^b ^c S. Goriely, J. José, M. Hernanz, M. Rayet, M. Arnould: He-detonation in sub-Chandrasekhar CO white dwarfs: A new insight into energetics and p-process nucleosynthesis. In: Astronomy and Astrophysics. Vol. 383, 2002, pp. L27-L30. doi: 10.1051 / 0004-6361: 20020088 .
↑ ^a ^b ^c C. Fröhlich, G. Martínez-Pinedo, M. Liebendörfer, F.-K. Thielemann, E. Bravo, WR Hix, K. Langanke, NT Zinner: Neutrino-Induced Nucleosynthesis of A> 64 Nuclei: The νp Process. In: Physical Review Letters. Vol. 96, 2006, article 142502. doi: 10.1103 / PhysRevLett.96.142502 .
^ ^A ^b T. Rauscher, A. Heger, RD Hoffman, SE Woosley: Nucleosynthesis in Massive Stars with Improved Nuclear and Stellar Physics. In: The Astrophysical Journal. Vol. 576, 2002, pp. 323-348. doi: 10.1086 / 341728 .
↑ C. Fröhlich et al. a .: Composition of the Innermost Core-Collapse Supernova Ejecta. In: The Astrophysical Journal. Vol. 637, 2006, pp. 415-426. doi: 10.1086 / 498224 .
^ WM Howard, SB Meyer, SE Woosley: A new site for the astrophysical gamma process. In: The Astrophysical Journal Letters. Vol. 373, 1991, pp. L5-L8. doi: 10.1086 / 186038 .
^ SE Woosley et al. a .: Models for Type I X-Ray Bursts with Improved Nuclear Physics. In: The Astrophysical Journal Supplement. Vol. 151, 2004, pp. 75-102. doi: 10.1086 / 381533 .

[b2fh-1] A ^b E. M. Burbidge , GR Burbidge , WA Fowler , Fred Hoyle : Synthesis of the Elements in Stars . In: Reviews of Modern Physics . 29, No. 4, 1957, pp. 547-650. doi : 10.1103 / RevModPhys.29.547 .

[cameron-2] A ^b A. GW Cameron : Nuclear Reactions in Stars and Nucleogenesis. In: Publications of the Astronomical Society of the Pacific. Vol. 69, 1957, pp. 201-222, bibcode : 1957PASP ... 69..201C .

[s-contrib-3] C. Arlandini, F. Käppeler, K. Wisshak, R. Gallino, M. Lugaro, M. Busso, O. Straniero: Neutron Capture in Low-Mass Asymptotic Giant Branch Stars: Cross Sections and Abundance Signatures. In: The Astrophysical Journal . Vol. 525, 1999, pp. 886-900. doi: 10.1086 / 307938 .

[r-contrib-4] Zs. Nemeth, F. Käppeler, C. Theis, T. Belgya, SW Yates: Nucleosynthesis in the Cd-In-Sn region. In: The Astrophysical Journal. Vol. 426, 1994, pp. 357-365. doi: 10.1086 / 174071 .

[dauphas-5] N. Dauphas, T. Rauscher, B. Marty, L. Reisberg: Short-lived p-nuclides in the early solar system and implications on the nucleosynthetic role of X-ray binaries. In: Nuclear Physics. Vol. A719, 2003, pp. C287-C295, doi: 10.1016 / S0375-9474 (03) 00934-5 , arxiv : astro-ph / 0211452 .

[arnould-6] ↑ ^a ^b ^c ^d ^e ^f ^g ^h ⁱ ^j M. Arnould, S. Goriely: The p-process of stellar nucleosynthesis: astrophysics and nuclear physics status. In: Physics Reports. 384, 2003, pp. 1-84.

[posnic2010-7] T. Rauscher: Origin of p-Nuclei in Explosive Nucleosynthesis. In: Proceedings of Science. arxiv : 1012.2213v1 .

[woosley-8] A ^b ^c S. E. Woosley, WM Howard: The p-process in supernovae. In: The Astrophysical Journal Supplement. Vol. 36, 1978, pp. 285-304. doi: 10.1086 / 190501 .

[nu-process-9] A ^b S. E. Woosley, DH Hartmann, RD Hoffman, WC Haxton: The ν-process. In: The Astrophysical Journal. Vol. 356, 1990, pp. 272-301. doi: 10.1086 / 168839 .

[schatz-10] H. Schatz u. a .: rp-Process Nucleosynthesis at Extreme Temperature and Density Conditions. In: Physics Reports. Vol. 294, 1998, 167-263. doi: 10.1016 / S0370-1573 (97) 00048-3 .

[rp-end-11] H. Schatz u. a .: End Point of the rp Process on Accreting Neutron Stars. In: Physical Review Letters. Vol. 86, 2001, pp. 3471-3474. doi: 10.1103 / PhysRevLett.86.3471 .

[pn-proc-12] S. Goriely, J. José, M. Hernanz, M. Rayet, M. Arnould: He-detonation in sub-Chandrasekhar CO white dwarfs: A new insight into energetics and p-process nucleosynthesis. In: Astronomy and Astrophysics. Vol. 383, 2002, pp. L27-L30. doi: 10.1051 / 0004-6361: 20020088 .

[nup-process-13] C. Fröhlich, G. Martínez-Pinedo, M. Liebendörfer, F.-K. Thielemann, E. Bravo, WR Hix, K. Langanke, NT Zinner: Neutrino-Induced Nucleosynthesis of A> 64 Nuclei: The νp Process. In: Physical Review Letters. Vol. 96, 2006, article 142502. doi: 10.1103 / PhysRevLett.96.142502 .

[rhhw-14] A ^b T. Rauscher, A. Heger, RD Hoffman, SE Woosley: Nucleosynthesis in Massive Stars with Improved Nuclear and Stellar Physics. In: The Astrophysical Journal. Vol. 576, 2002, pp. 323-348. doi: 10.1086 / 341728 .

[nup-process2-15] C. Fröhlich et al. a .: Composition of the Innermost Core-Collapse Supernova Ejecta. In: The Astrophysical Journal. Vol. 637, 2006, pp. 415-426. doi: 10.1086 / 498224 .

[Ia-gamma-16] WM Howard, SB Meyer, SE Woosley: A new site for the astrophysical gamma process. In: The Astrophysical Journal Letters. Vol. 373, 1991, pp. L5-L8. doi: 10.1086 / 186038 .

[x-bursts-17] SE Woosley et al. a .: Models for Type I X-Ray Bursts with Improved Nuclear Physics. In: The Astrophysical Journal Supplement. Vol. 151, 2004, pp. 75-102. doi: 10.1086 / 381533 .