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.


Part of the Chart of Nuclides showing some stable or nearly-stable s-, r-, and p-nuclei

The classical, ground-breaking works of Burbidge, Burbidge, Fowler and Hoyle (1957)[1] and of A. G. W. Cameron (1957)[2] showed how the majority of naturally occurring nuclides beyond the element iron can be made in two kinds of neutron capture processes, the s- and the r-process. Some proton-rich nuclides found in nature are not reached in these processes and therefore at least one additional process is required to synthesize them. These nuclei are called p-nuclei.

Since the definition of the p-nuclei depends on the current knowledge of the s- and r-process (see also nucleosynthesis), the original list of 35 p-nuclei may be modified over the years, as indicated in the Table below. For example, it is recognized today that the abundances of 152Gd and 164Er contain at least strong contributions from the s-process.[3] This also seems to apply to those of 113In and 115Sn, which additionally could be made in the r-process in small amounts.[4]

The long-lived radionuclides 92Nb, 97Tc, 98Tc and 146Sm are not among the classically defined p-nuclei as they no longer occur naturally on Earth. By the above definition, however, they are also p-nuclei because they cannot be made in either the s- or the r-process. From the discovery of their decay products in presolar grains it can be inferred that at least 92Nb and 146Sm were present in the solar nebula. This offers the possibility to estimate the time since the last production of these p-nuclei before the formation of the solar system.[5]

p-nuclei are very rare. Those isotopes of an element which are p-nuclei are less abundant typically by factors of ten to one thousand than the other isotopes of the same element. The abundances of p-nuclei can only be determined in geochemical investigations and by analysis of meteoritic material and presolar grains. They cannot be identified in stellar spectra. Therefore, the knowledge of p-abundances is restricted to those of the Solar System and it is unknown whether the solar abundances of p-nuclei are typical for the Milky Way.[6]

List of p-nuclei
Nuclide Comment
78Kr long-lived radionuclide
92Nb long-lived radionuclide; not a classical p-nucleus but cannot be made in s- and r-processes
97Tc long-lived radionuclide; not a classical p-nucleus but cannot be made in s- and r-processes
98Tc long-lived radionuclide; not a classical p-nucleus but cannot be made in s- and r-processes
113In (partially) made in the s-process? Contributions from the r-process?
115Sn (partially) made in the s-process? Contributions from the r-process?
130Ba long-lived radionuclide
138La long-lived radionuclide; made in the ν-process
146Sm long-lived radionuclide; not a classical p-nucleus but cannot be made in s- and r-processes
152Gd long-lived radionuclide; (partially) made in the s-process?
164Er (partially) made in the s-process?
174Hf long-lived radionuclide
180mTa (partially) made in the ν-process; contributions from the s-process?
180W long-lived radionuclide
190Pt long-lived radionuclide

Origin of the p-nuclei

The astrophysical production of p-nuclei is not completely understood yet. The favored γ-process (see below) in core-collapse supernovae cannot produce all p-nuclei in sufficient amounts, according to current computer simulations. This is why additional production mechanisms and astrophysical sites are under investigation, as outlined below. It is also conceivable that there is not just a single process responsible for all p-nuclei but that different processes in a number of astrophysical sites produce certain ranges of p-nuclei.[7]

In the search for the relevant processes creating p-nuclei, the usual way is to identify the possible production mechanisms (processes) and then to investigate their possible realization in various astrophysical sites. The same logic is applied in the discussion below.

Basics of p-nuclide production

In principle, there are two ways to produce proton-rich nuclides: by successively adding protons to a nuclide (these are nuclear reactions of type (p,γ) or by removing neutrons from a nucleus through sequences of photodisintegrations of type (γ,n).[6][7]

Under conditions encountered in astrophysical environments it is difficult to obtain p-nuclei through proton captures because the Coulomb barrier of a nucleus increases with increasing proton number. A proton requires more energy to be incorporated (captured) into an atomic nucleus when the Coulomb barrier is higher. The available average energy of the protons is determined by the temperature of the stellar plasma. Increasing the temperature, however, also speeds up the (γ,p) photodisintegrations which counteract the (p,γ) captures. The only alternative avoiding this would be to have a very large number of protons available so that the effective number of captures per second is large even at low temperature. In extreme cases (as discussed below) this leads to the synthesis of extremely short-lived radionuclides which decay to stable nuclides only after the captures cease.[6][7]

Appropriate combinations of temperature and proton density of a stellar plasma have to be explored in the search of possible production mechanisms for p-nuclei. Further parameters are the time available for the nuclear processes, and number and type of initially present nuclides (seed nuclei).

Possible processes

The p-process

In a p-process it is suggested that p-nuclei were made through a few proton captures on stable nuclides. The seed nuclei originate from the s- and r-process and are already present in the stellar plasma. As outlined above, there are serious difficulties explaining all p-nuclei through such a process although it was originally suggested to achieve exactly this.[1][2][6] It was shown later that the required conditions are not reached in stars or stellar explosions.[8]

Based on its historical meaning, the term p-process is sometimes sloppily used for any process synthesizing p-nuclei, even when no proton captures are involved.

The γ-process

p-Nuclei can also be obtained by photodisintegration of s- and r-process nuclei. At temperatures around 2–3 gigakelvins (GK) and short process time of a few seconds (this requires an explosive process) photodisintegration of the pre-existing nuclei will remain small, just enough to produce the required tiny abundances of p-nuclei.[6][9] This is called γ-process because the photodisintegration proceeds by nuclear reactions of the types (γ,n), (γ,α) and (γ,p), which are caused by highly energetic photons (Gamma rays).[9]

The ν-Process

If a sufficiently intensive source of neutrinos is available, nuclear reactions can directly produce certain nuclides, for example 7Li, 11B, 19F, 138La in core-collapse supernovae.[10]

Rapid proton capture processes

In a p-process protons are added to stable or weakly radioactive atomic nuclei. If there is a high proton density in the stellar plasma, even short-lived radionuclides can capture one or more protons before they beta decay. This quickly moves the nucleosynthesis path from the region of stable nuclei to the very proton-rich side of the Chart of Nuclides. This is called rapid proton-capture.[7]

Here, a series of (p,γ) reactions proceeds until either the beta decay of a nucleus is faster than a further proton capture, or the proton drip line is reached. Both cases lead to one or several sequential beta decays until a nucleus is produced which again can capture protons before it beta decays. Then the proton capture sequences continue.

It is possible to cover the region of the lightest nuclei up to 56Ni within a second because both proton captures and beta decays are fast. Starting with 56Ni, however, a number of waiting points are encountered in the reaction path. These are nuclides which both have relatively long half-lives (compared to the process timescale) and can only slowly add another proton (that is, their cross section for (p,γ) reactions is small). Examples for such waiting points are: 56Ni, 60Zn, 64Ge, 68Se. Further waiting points may be important, depending on the detailed conditions and location of the reaction path. It is typical for such waiting points to show half-lives of minutes to days. Thus, they considerably increase the time required to continue the reaction sequences. If the conditions required for this rapid proton capture are only present for a short time (the timescale of explosive astrophysical events is of the order of seconds), the waiting points limit or hamper the continuation of the reactions to heavier nuclei.[11]

In order to produce p-nuclei, the process path has to encompass nuclides bearing the same mass number (but usually containing more protons) as the desired p-nuclei. These nuclides are then converted into p-nuclei through sequences of beta decays after the rapid proton captures ceased.

Variations of the main category rapid proton captures are the rp-, pn-, and νp-processes, which will be briefly outlined below.

The so-called rp-process (rp is for rapid proton capture) is the purest form of the rapid proton capture process described above. At proton densities of more than 1028 protons/cm3 and temperatures around 2 GK the reaction path is close to the proton drip line.[11] The waiting points can be bridged provided that the process time is 10-600 s. Waiting-point nuclides are produced with larger abundances while the production of nuclei "behind" each waiting-point is more and more suppressed.

A definitive endpoint is reached close to 107Te because the reaction path runs into a region of nuclides which decay preferably by alpha decay and thus loop the path back onto itself.[12] Therefore, an rp-process would only be able to produce p-nuclei with mass numbers less than or equal to 107.

The waiting points in rapid proton capture processes can be avoided by (n,p) reactions which are much faster than proton captures on or beta decays of waiting points nuclei. This results in a considerable reduction of the time required to build heavy elements and allows an efficient production within seconds.[6] This requires, however, a (small) supply of free neutrons which are usually not present in such proton-rich plasmas. One way to obtain them is to release them through other reactions occurring simultaneously as the rapid proton captures. This is called neutron-rich rapid proton capture or pn-process.[13]

Another possibility to obtain the neutrons required for the accelerating (n,p) reactions in proton-rich environments is to use the anti-neutrino capture on protons (

), turning a proton and an anti-neutrino into a positron and a neutron. Since (anti-)neutrinos interact only very weakly with protons, a high flux of anti-neutrinos has to act on a plasma with high proton density. This is called νp-process.[14]

Possible synthesis sites

Core-collapse supernovae

Massive stars end their life in a core-collapse supernova. In such a supernova, a shockfront from an explosion runs from the center of the star through its outer layers and ejects these. When the shockfront reaches the O/Ne-shell of the star (see also stellar evolution), the conditions for a γ-process are reached for 1-2 s.

Although the majority of p-nuclei can be made in this way, some mass regions of p-nuclei turn out to be problematic in model calculations. It has been known already for decades that p-nuclei with mass numbers A < 100 cannot be produced in a γ-process.[6][9] Modern simulations also show problems in the range 150 ≤ A ≤ 165.[7][15]

The p-nucleus 138La is not produced in the γ-process but it can be made in a ν-process. A hot neutron star is made in the center of such a core-collapse supernova and it radiates neutrinos with high intensity. The neutrinos interact also with the outer layers of the exploding star and cause nuclear reactions which create 138La, among other nuclei.[10][15] Also 180mTa may receive a contribution from this ν-process.

It was suggested[14] to supplement the γ-process in the outer layers of the star by another process, occurring in the deepest layers of the star, close to the neutron star but still being ejected instead of falling onto the neutron star surface. Due to the initially high flow of neutrinos from the forming neutron star, these layers become extremely proton-rich through the reaction

. Although the anti-neutrino flux is initially weaker a few neutrons will be created, nevertheless, because of the large number of protons. This allows a νp-process in these deep layers. Because of the short timescale of the explosion and the high Coulomb barrier of the heavier nuclei, such a νp-process could possibly only produce the lightest p-nuclei. Which nuclei are made and how much of them depends sensitively on many details in the simulations and also on the actual explosion mechanism of a core-collapse supernova, which still is not completely understood.[14][16]

Thermonuclear supernovae

A thermonuclear supernova is the explosion of a white dwarf in a binary star system, triggered by thermonuclear reactions in matter from a companion star accreted on the surface of the white dwarf. The accreted matter is rich in hydrogen (protons) and helium (α particles) and becomes hot enough to allow nuclear reactions.

A number of models for such explosions are discussed in literature, of which two were explored regarding the prospect of producing p-nuclei. None of these explosions release neutrinos, therefore rendering ν- and νp-process impossible. Conditions required for the rp-process are also not attained.

Details of the possible production of p-nuclei in such supernovae depend sensitively on the composition of the matter accreted from the companion star (the seed nuclei for all subsequent processes). Since this can change considerably from star to star, all statements and models of p-production in thermonuclear supernovae are prone to large uncertainties.[6]

The consensus model of thermonuclear supernovae postulates that the white dwarf explodes after exceeding the Chandrasekhar limit by the accretion of matter because the contraction and heating ignites explosive carbon burning under degenerate conditions. A nuclear burning front runs through the white dwarf from the inside out and tears it apart. Then the outermost layers closely beneath the surface of the white dwarf (containing 0.05 solar masses of matter) exhibit the right conditions for a γ-process.[17]

The p-nuclei are made in the same way as in the γ-process in core-collapse supernovae and also the same difficulties are encountered. In addition, 138La and 180mTa are not produced. A variation of the seed abundances by assuming increased s-process abundances only scales the abundances of the resulting p-nuclei without curing the problems of relative underproduction in the nuclear mass ranges given above.[6]

In a subclass of type Ia supernovae, the so-called subChandrasekhar supernova, the white dwarf may explode long before it reaches the Chandrasekhar limit because nuclear reactions in the accreted matter can already heat the white dwarf during its accretion phase and trigger explosive carbon burning prematurely. Helium-rich accretion favors this type of explosion. Helium burning ignites degeneratively on the bottom of the accreted helium layer and causes two shockfronts. The one running inwards ignites the carbon explosion. The outwards moving front heats the outer layers of the white dwarf and ejects them. Again, these outer layers are site to a γ-process at temperatures of 2-3 GK. Due to the presence of α particles (helium nuclei), however, additional nuclear reactions become possible. Among those are such which release a large number of neutrons, such as 18O(α,n)21Ne, 22Ne(α,n)25Mg, and 26Mg(α,n)29Si. This allows a pn-process in that part of the outer layers which experiences temperatures above 3 GK.[6][13]

Those light p-nuclei which are underproduced in the γ-process can be so efficiently made in the pn-process that they even show much larger abundances than the other p-nuclei. To obtain the observed solar relative abundances, a strongly enhanced s-process seed (by factors of 100-1000 or more) has to be assumed which increases the yield of heavy p-nuclei from the γ-process.[6][13]

Neutron stars in binary star systems

A neutron star in a binary star system can also accrete matter from the companion star on its surface. Combined hydrogen and helium burning ignites when the accreted layer of degenerate matter reaches a density of 105106 g/cm3 and a temperature exceeding 0.2 GK. This leads to thermonuclear burning comparable to what happens in the outwards moving shockfront of subChandrasekhar supernovae. The neutron star itself is not affected by the explosion and therefore the nuclear reactions in the accreted layer can proceed longer than in an explosion. This allows to establish an rp-process. It will continue until either all free protons are used up or the burning layer has expanded due to the increase in temperature and its density falls below the one required for the nuclear reactions.[11]

It was shown that the properties of X-ray bursts in the Milky Way can be explained by an rp-process on the surface of accreting neutron stars.[18] It remains unclear, yet, whether matter (and if, how much matter) can be ejected and escape the gravitational field of the neutron star. Only if this is the case can such objects be considered as possible sources of p-nuclei. Even if this is corroborated, the demonstrated endpoint of the rp-process limits the production to the light p-nuclei (which are underproduced in core-collapse supernovae).[12]

See also


  1. ^ a b E. M. Burbidge; G. R. Burbidge; W. A. Fowler; Fred Hoyle (1957). "Synthesis of the Elements in Stars" (PDF). Reviews of Modern Physics. 29 (4): 547–650. Bibcode:1957RvMP...29..547B. doi:10.1103/RevModPhys.29.547.
  2. ^ a b A. G. W. Cameron: Nuclear Reactions in Stars and Nucleogenesis. In: Publications of the Astronomical Society of the Pacific, Vol. 69, 1957, p. 201-222. (online)
  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, p. 886-900. ( doi:10.1086/307938)
  4. ^ Zs. Nemeth, F. Käppeler, C. Theis, T. Belgya, S. W. Yates: Nucleosynthesis in the Cd-In-Sn region. In: The Astrophysical Journal, Vol. 426, 1994, p. 357-365. ( doi:10.1086/174071)
  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, p. C287-C295 ( doi:10.1016/S0375-9474(03)00934-5, arXiv.org:astro-ph/0211452)
  6. ^ a b c d e f g h i j k M. Arnould, S. Goriely: The p-process of stellar nucleosynthesis: astrophysics and nuclear physics status. In: Physics Reports 384, 2003, p. 1-84.
  7. ^ a b c d e T. Rauscher: Origin of p-Nuclei in Explosive Nucleosynthesis. In: Proceedings of Science XI_059.pdf PoS(NIC XI)059, 2010 (arXiv.org:1012.2213)
  8. ^ J. Audouze, J. W. Truran: P-process nucleosynthesis in postshock supernova envelope environments. In: The Astrophysical Journal, Vol. 202, 1975, p. 204-213. ( doi:10.1086/153965)
  9. ^ a b c S. E. Woosley, W. M. Howard: The p-process in supernovae. In: The Astrophysical Journal Supplement, Vol. 36, 1978, p. 285–304. (doi:10.1086/190501)
  10. ^ a b S. E. Woosley, D. H. Hartmann, R. D. Hoffman, W. C. Haxton: The ν-process. In: The Astrophysical Journal, Vol. 356, 1990, p. 272-301. ( doi:10.1086/168839)
  11. ^ a b c H. Schatz, et al.: rp-Process Nucleosynthesis at Extreme Temperature and Density Conditions. In: Physics Reports, Vol. 294, 1998, p. 167-263. ( doi:10.1016/S0370-1573(97)00048-3)
  12. ^ a b H. Schatz, et al.: End Point of the rp Process on Accreting Neutron Stars. In: Physical Review Letters, Vol. 86, 2001, p. 3471-3474. ([1] doi:10.1016/10.1103/PhysRevLett.86.3471)
  13. ^ 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, p. L27-L30. ( doi:10.1051/0004-6361:20020088)
  14. ^ a b c C. Fröhlich, G. Martínez-Pinedo, M. Liebendörfer, F.-K. Thielemann, E. Bravo, W. R. Hix, K. Langanke, N. T. 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)
  15. ^ a b T. Rauscher, A. Heger, R. D. Hoffman, S. E. Woosley: Nucleosynthesis in Massive Stars with Improved Nuclear and Stellar Physics. In: The Astrophysical Journal, Vol. 576, 2002, p. 323-348. ( doi:10.1086/341728)
  16. ^ C. Fröhlich, et al.: Composition of the Innermost Core-Collapse Supernova Ejecta. In: The Astrophysical Journal, Vol. 637, 2006, p. 415-426. ( doi:10.1086/498224)
  17. ^ W. M. Howard, S. B. Meyer, S. E. Woosley: A new site for the astrophysical gamma-process. In: The Astrophysical Journal Letters, Vol. 373, 1991, p. L5-L8. ( doi:10.1086/186038)
  18. ^ S. E. Woosley, et al.: Models for Type I X-Ray Bursts with Improved Nuclear Physics. In: The Astrophysical Journal Supplement, Vol. 151, 2004, p. 75-102. ( doi:10.1086/381553)

Cerium is a chemical element with the symbol Ce and atomic number 58. Cerium is a soft, ductile and silvery-white metal that tarnishes when exposed to air, and it is soft enough to be cut with a knife. Cerium is the second element in the lanthanide series, and while it often shows the +3 oxidation state characteristic of the series, it also exceptionally has a stable +4 state that does not oxidize water. It is also considered one of the rare-earth elements. Cerium has no biological role and is not very toxic.

Despite always occurring in combination with the other rare-earth elements in minerals such as those of the monazite and bastnäsite groups, cerium is easy to extract from its ores, as it can be distinguished among the lanthanides by its unique ability to be oxidized to the +4 state. It is the most common of the lanthanides, followed by neodymium, lanthanum, and praseodymium. It is the 26th-most abundant element, making up 66 ppm of the Earth's crust, half as much as chlorine and five times as much as lead.

Cerium was the first of the lanthanides to be discovered, in Bastnäs, Sweden by Jöns Jakob Berzelius and Wilhelm Hisinger in 1803, and independently by Martin Heinrich Klaproth in Germany in the same year. In 1839 Carl Gustaf Mosander became the first to isolate the metal. Today, cerium and its compounds have a variety of uses: for example, cerium(IV) oxide is used to polish glass and is an important part of catalytic converters. Cerium metal is used in ferrocerium lighters for its pyrophoric properties. Cerium-doped YAG phosphor is used in conjunction with blue light-emitting diodes to produce white light in most commercial white LED light sources.

Index of physics articles (P)

The index of physics articles is split into multiple pages due to its size.

To navigate by individual letter use the table of contents below.


Indium is a chemical element with the symbol In and atomic number 49. Indium is the softest metal that is not considered to be an alkali metal. It is a silvery-white metal that resembles Tin(Sn) in appearance. It is a post-transition metal that makes up 0.21 parts per million of the Earth's crust. Indium has a melting point higher than sodium and gallium, but lower than lithium and tin. Chemically, indium is similar to gallium and thallium, and it is largely intermediate between the two in terms of its properties. Indium was discovered in 1863 by Ferdinand Reich and Hieronymous Theodor Richter by spectroscopic methods. They named it for the indigo blue line in its spectrum. Indium was isolated the next year.

Indium is a minor component in zinc sulfide ores and is produced as a byproduct of zinc refinement. It is most notably used in the semiconductor industry, in low-melting-point metal alloys such as solders, in soft-metal high-vacuum seals, and in the production of transparent conductive coatings of indium tin oxide (ITO) on glass. Indium is considered a technology-critical element.

Indium has no biological role, though its compounds are somewhat toxic when injected into the bloodstream. Most occupational exposure is through ingestion, from which indium compounds are not absorbed well, and inhalation, from which they are moderately absorbed.


Lanthanum is a chemical element with the symbol La and atomic number 57. It is a soft, ductile, silvery-white metal that tarnishes slowly when exposed to air and is soft enough to be cut with a knife. It is the eponym of the lanthanide series, a group of 15 similar elements between lanthanum and lutetium in the periodic table, of which lanthanum is the first and the prototype. It is also sometimes considered the first element of the 6th-period transition metals, which would put it in group 3, although lutetium is sometimes placed in this position instead, for example in the left-step model. Lanthanum is traditionally counted among the rare earth elements. The usual oxidation state is +3. Lanthanum has no biological role in humans but is essential to some bacteria. It is not particularly toxic to humans but does show some antimicrobial activity.

Lanthanum usually occurs together with cerium and the other rare earth elements. Lanthanum was first found by the Swedish chemist Carl Gustav Mosander in 1839 as an impurity in cerium nitrate – hence the name lanthanum, from the Ancient Greek λανθάνειν (lanthanein), meaning "to lie hidden". Although it is classified as a rare earth element, lanthanum is the 28th most abundant element in the Earth's crust, almost three times as abundant as lead. In minerals such as monazite and bastnäsite, lanthanum composes about a quarter of the lanthanide content. It is extracted from those minerals by a process of such complexity that pure lanthanum metal was not isolated until 1923.

Lanthanum compounds have numerous applications as catalysts, additives in glass, carbon arc lamps for studio lights and projectors, ignition elements in lighters and torches, electron cathodes, scintillators, GTAW electrodes, and other things. Lanthanum carbonate is used as a phosphate binder in cases of renal failure.

List of unsolved problems in astronomy

Some of the unsolved problems in astronomy are theoretical, meaning that existing theories seem incapable of explaining a certain observed phenomenon or experimental result. The others are experimental, meaning that there is a difficulty in creating an experiment to test a proposed theory or investigate a phenomenon in greater detail. Some unresolved questions in astronomy pertain to one-off events, unusual occurrences that have not repeated and whose causes therefore remain unclear.

List of unsolved problems in physics

Some of the major unsolved problems in physics are theoretical, meaning that existing theories seem incapable of explaining a certain observed phenomenon or experimental result. The others are experimental, meaning that there is a difficulty in creating an experiment to test a proposed theory or investigate a phenomenon in greater detail.

There are still some deficiencies in the Standard Model of physics, such as the origin of mass, the strong CP problem, neutrino mass, matter–antimatter asymmetry, and the nature of dark matter and dark energy. Another problem lies within the mathematical framework of the Standard Model itself—the Standard Model is inconsistent with that of general relativity, to the point that one or both theories break down under certain conditions (for example within known spacetime singularities like the Big Bang and the centers of black holes beyond the event horizon).


The term p-process (p is for proton) is used in two ways in the scientific literature concerning the astrophysical origin of the elements (nucleosynthesis). Originally it referred to a proton capture process which is the source of certain, naturally occurring, proton-rich isotopes of the elements from selenium to mercury. These nuclides are called p-nuclei and their origin is still not completely understood. Although it was shown that the originally suggested process cannot produce the p-nuclei, later on the term p-process was sometimes used to generally refer to any nucleosynthesis process supposed to be responsible for the p-nuclei.Often, the two meanings are confused. Recent scientific literature therefore suggests to use the term p-process only for the actual proton capture process, as it is customary with other nucleosynthesis processes in astrophysics.

Proton capture

Proton capture is a nuclear reaction in which an atomic nucleus and one or more protons collide and merge to form a heavier nucleus.

Since protons have positive electric charge, they are repelled electrostatically by the positively charged nucleus. Therefore, it is more difficult for protons to enter the nucleus compared to neutrally charged neutrons .

Proton capture plays an important role in the cosmic nucleosynthesis of proton rich isotopes. In stars it can proceed in two ways: as a rapid (rp-process) or a slow process (p-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.


A supernova ( plural: supernovae or supernovas, abbreviations: SN and SNe) is a transient astronomical event that occurs during the last stellar evolutionary stages of the life of a massive star, whose dramatic and catastrophic destruction is marked by one final, titanic explosion. This causes the sudden appearance of a "new" bright star, before slowly fading from sight over several weeks or months or years.

Supernovae are more energetic than novae. In Latin, nova means "new", referring astronomically to what appears to be a temporary new bright star. Adding the prefix "super-" distinguishes supernovae from ordinary novae, which are far less luminous. The word supernova was coined by Walter Baade and Fritz Zwicky in 1931.Only three Milky Way, naked-eye supernova events have been observed during the last thousand years, though many have been seen in other galaxies. The most recent directly observed supernova in the Milky Way was Kepler's Supernova in 1604, but two more recent supernova remnants have also been found. Statistical observations of supernovae in other galaxies suggest they occur on average about three times every century in the Milky Way, and that any galactic supernova would almost certainly be observable with modern astronomical telescopes.

Supernovae may expel much, if not all, of the material away from a star at velocities up to 30,000 km/s or 10% of the speed of light. This drives an expanding and fast-moving shock wave into the surrounding interstellar medium, and in turn, sweeping up an expanding shell of gas and dust, which is observed as a supernova remnant. Supernovae create, fuse and eject the bulk of the chemical elements produced by nucleosynthesis. Supernovae play a significant role in enriching the interstellar medium with the heavier atomic mass chemical elements. Furthermore, the expanding shock waves from supernovae can trigger the formation of new stars. Supernova remnants are expected to accelerate a large fraction of galactic primary cosmic rays, but direct evidence for cosmic ray production was found only in a few of them so far. They are also potentially strong galactic sources of gravitational waves.Theoretical studies indicate that most supernovae are triggered by one of two basic mechanisms: the sudden re-ignition of nuclear fusion in a degenerate star or the sudden gravitational collapse of a massive star's core. In the first instance, a degenerate white dwarf may accumulate sufficient material from a binary companion, either through accretion or via a merger, to raise its core temperature enough to trigger runaway nuclear fusion, completely disrupting the star. In the second case, the core of a massive star may undergo sudden gravitational collapse, releasing gravitational potential energy as a supernova. While some observed supernovae are more complex than these two simplified theories, the astrophysical collapse mechanics have been established and accepted by most astronomers for some time.

Owing to the wide range of astrophysical consequences of these events, astronomers now deem supernova research, across the fields of stellar and galactic evolution, as an especially important area for investigation.


Tin is a chemical element with the symbol Sn (from Latin: stannum) and atomic number 50. Tin is a silvery white metal that characteristicly has a faint yellow hue due to slight oxidation. Tin, like indium, is soft enough to be cut without much force. When a bar of tin is bent the so-called "tin cry" can be heard as a result of sliding tin crystals reforming; this trait is shared by indium, cadmium and frozen mercury. Pure tin after solidifying keeps a mirror-like appearance similar to most metals, however most Tin alloys such as in Pewter, the metal soldifies with a dull gray color. Tin is a post-transition metal in group 14 of the periodic table of elements. It is obtained chiefly from the mineral cassiterite, which contains stannic oxide, SnO2. Tin shows a chemical similarity to both of its neighbors in group 14, germanium and lead, and has two main oxidation states, +2 and the slightly more stable +4. Tin is the 49th most abundant element and has, with 10 stable isotopes, the largest number of stable isotopes in the periodic table, thanks to its magic number of protons. It has two main allotropes: at room temperature, the stable allotrope is β-tin, a silvery-white, malleable metal, but at low temperatures, it transforms into the less dense grey α-tin, which has the diamond cubic structure. Metallic tin does not easily oxidize in air.

The first tin alloy used on a large scale was bronze, made of 1/8 tin and 7/8 copper, from as early as 3000 BC. After 600 BC, pure metallic tin was produced. Pewter, which is an alloy of 85–90% tin with the remainder commonly consisting of copper, antimony, and lead, was used for flatware from the Bronze Age until the 20th century. In modern times, tin is used in many alloys, most notably tin/lead soft solders, which are typically 60% or more tin, and in the manufacture of transparent, electrically conducting films of indium tin oxide in optoelectronic applications. Another large application for tin is corrosion-resistant tin plating of steel. Because of the low toxicity of inorganic tin, tin-plated steel is widely used for food packaging as tin cans. However, some organotin compounds can be almost as toxic as cyanide.

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