Stellar nucleosynthesis

Stellar nucleosynthesis is the theory explaining the creation (nucleosynthesis) of chemical elements by nuclear fusion reactions between atoms within stars. Stellar nucleosynthesis has occurred continuously since the original creation of hydrogen, helium and lithium during the Big Bang. It is a highly predictive theory that today yields excellent agreement between calculations based upon it and the observed abundances of the elements. It explains why the observed abundances of elements in the universe grow over time and why some elements and their isotopes are much more abundant than others. The theory was initially proposed by Fred Hoyle in 1946,[1] who later refined it in 1954.[2] Further advances were made, especially to nucleosynthesis by neutron capture of the elements heavier than iron, by Margaret Burbidge, Geoffrey Burbidge, William Alfred Fowler and Hoyle in their famous 1957 B2FH paper,[3] which became one of the most heavily cited papers in astrophysics history.

Stars evolve because of changes in their composition (the abundance of their constituent elements) over their lifespans, first by burning hydrogen (main sequence star), then helium (red giant star), and progressively burning higher elements. However, this does not by itself significantly alter the abundances of elements in the universe as the elements are contained within the star. Later in its life, a low-mass star will slowly eject its atmosphere via stellar wind, forming a planetary nebula, while a higher–mass star will eject mass via a sudden catastrophic event called a supernova. The term supernova nucleosynthesis is used to describe the creation of elements during the evolution and explosion of a pre-supernova massive star (12–35 times the mass of the sun). Those massive stars are the most prolific source of new isotopes from carbon (Z = 6) to nickel (Z = 28).

The advanced sequence of burning fuels is driven by gravitational collapse and its associated heating, resulting in the subsequent burning of carbon, oxygen and silicon. However, most of the nucleosynthesis in the mass range A = 28–56 (from silicon to nickel) is actually caused by the upper layers of the star collapsing onto the core, creating a compressional shock wave rebounding outward. The shock front briefly raises temperatures by roughly 50%, thereby causing furious burning for about a second. This final burning in massive stars, called explosive nucleosynthesis or supernova nucleosynthesis, is the final epoch of stellar nucleosynthesis.

A stimulus to the development of the theory of nucleosynthesis was the discovery of variations in the abundances of elements found in the universe. The need for a physical description was already inspired by the relative abundances of isotopes of the chemical elements in the solar system. Those abundances, when plotted on a graph as a function of atomic number of the element, have a jagged sawtooth shape that varies by factors of tens of millions (see history of nucleosynthesis theory).[4] This suggested a natural process that is not random. A second stimulus to understanding the processes of stellar nucleosynthesis occurred during the 20th century, when it was realized that the energy released from nuclear fusion reactions accounted for the longevity of the Sun as a source of heat and light.[5]

History

Arthur Stanley Eddington
In 1920, Arthur Eddington proposed that stars obtained their energy from nuclear fusion of hydrogen to form helium and also raised the possibility that the heavier elements are produced in stars.

In 1920, Arthur Eddington, on the basis of the precise measurements of atomic masses by F.W. Aston and a preliminary suggestion by Jean Perrin, proposed that stars obtained their energy from nuclear fusion of hydrogen to form helium and raised the possibility that the heavier elements are produced in stars.[6][7][8] This was a preliminary step toward the idea of nucleosynthesis. In 1928, George Gamow derived what is now called the Gamow factor, a quantum-mechanical formula that gave the probability of bringing two nuclei sufficiently close for the strong nuclear force to overcome the Coulomb barrier. The Gamow factor was used in the decade that followed by Atkinson and Houtermans and later by Gamow himself and Edward Teller to derive the rate at which nuclear reactions would proceed at the high temperatures believed to exist in stellar interiors.

In 1939, in a paper entitled "Energy Production in Stars", Hans Bethe analyzed the different possibilities for reactions by which hydrogen is fused into helium.[9] He defined two processes that he believed to be the sources of energy in stars. The first one, the proton–proton chain reaction, is the dominant energy source in stars with masses up to about the mass of the Sun. The second process, the carbon–nitrogen–oxygen cycle, which was also considered by Carl Friedrich von Weizsäcker in 1938, is more important in more massive main-sequence stars.[10] These works concerned the energy generation capable of keeping stars hot. A clear physical description of the proton–proton chain and of the CNO cycle appears in a 1968 textbook[5]. Bethe's two papers did not address the creation of heavier nuclei, however. That theory was begun by Fred Hoyle in 1946 with his argument that a collection of very hot nuclei would assemble thermodynamically into iron[1]. Hoyle followed that in 1954 with a large paper describing how advanced fusion stages within massive stars would synthesize the elements from carbon to iron in mass[2]. This is the first work of stellar nucleosynthesis[11]. It and Hoyle's 1954 paper provided the roadmap to how the most abundant elements on Earth had been synthesized within stars from their initial hydrogen and helium, making clear how those abundant elements increased their galactic abundances as the galaxy aged.

Hoyle's theory was expanded to other processes, beginning with the publication of a review paper in 1957 by Burbidge, Burbidge, Fowler and Hoyle (commonly referred to as the B2FH paper).[3] This review paper collected and refined earlier research into a heavily cited picture that gave promise of accounting for the observed relative abundances of the elements; but it did not itself enlarge Hoyle's 1954 picture for the origin of primary nuclei as much as many assumed, except in the understanding of nucleosynthesis of those elements heavier than iron by neutron capture. Significant improvements were made by Alastair G. W. Cameron and by Donald D. Clayton. Cameron presented his own independent approach[12] (following Hoyle's approach for the most part) of nucleosynthesis. He introduced computers into time-dependent calculations of evolution of nuclear systems. Clayton calculated the first time-dependent models of the S-process[13] and of the R-process,[14] as well as of the burning of silicon into the abundant alpha-particle nuclei and iron-group elements,[15][16] and discovered radiogenic chronologies[17] for determining the age of the elements. The entire research field expanded rapidly in the 1970s.

Nucleosynthesis in a star
Cross section of a supergiant showing nucleosynthesis and elements formed.

Key reactions

Nucleosynthesis periodic table
A version of the periodic table indicating the origins – including stellar nucleosynthesis – of the elements. Elements above 94 are manmade and are not included.

The most important reactions in stellar nucleosynthesis:

Hydrogen fusion

FusionintheSun
Proton–proton chain reaction
CNO Cycle
CNO-I cycle
The helium nucleus is released at the top-left step.

Hydrogen fusion (nuclear fusion of four protons to form a helium-4 nucleus[18]) is the dominant process that generates energy in the cores of main-sequence stars. It is also called "hydrogen burning", which should not be confused with the chemical combustion of hydrogen in an oxidizing atmosphere. There are two predominant processes by which stellar hydrogen fusion occurs: proton-proton chain and the carbon-nitrogen-oxygen (CNO) cycle. Ninety percent of all stars, with the exception of white dwarfs, are fusing hydrogen by these two processes.

In the cores of lower-mass main-sequence stars such as the Sun, the dominant energy production process is the proton–proton chain reaction. This creates a helium-4 nucleus through a sequence of chain reactions that begin with the fusion of two protons to form a deuterium nucleus (one proton plus one neutron) along with an ejected positron and neutrino.[19] In each complete fusion cycle, the proton–proton chain reaction releases about 26.2 MeV.[19] The proton–proton chain reaction cycle is relatively insensitive to temperature; a 10% rise of temperature would increase energy production by this method by 46%, hence, this hydrogen fusion process can occur in up to a third of the star's radius and occupy half the star's mass. For stars above 35% of the Sun's mass,[20] the energy flux toward the surface is sufficiently low and energy transfer from the core region remains by radiative heat transfer, rather than by convective heat transfer.[21] As a result, there is little mixing of fresh hydrogen into the core or fusion products outward.

In higher-mass stars, the dominant energy production process is the CNO cycle, which is a catalytic cycle that uses nuclei of carbon, nitrogen and oxygen as intermediaries and in the end produces a helium nucleus as with the proton-proton chain.[19] During a complete CNO cycle, 25.0 MeV of energy is released. The difference in energy production of this cycle, compared to the proton–proton chain reaction, is accounted for by the energy lost through neutrino emission.[19] The CNO cycle is very temperature sensitive, a 10% rise of temperature would produce a 350% rise in energy production. About 90% of the CNO cycle energy generation occurs within the inner 15% of the star's mass, hence it is strongly concentrated at the core.[22] This results in such an intense outward energy flux that convective energy transfer become more important than does radiative transfer. As a result, the core region becomes a convection zone, which stirs the hydrogen fusion region and keeps it well mixed with the surrounding proton-rich region.[23] This core convection occurs in stars where the CNO cycle contributes more than 20% of the total energy. As the star ages and the core temperature increases, the region occupied by the convection zone slowly shrinks from 20% of the mass down to the inner 8% of the mass.[22] Our Sun produces 10% of its energy from the CNO cycle.

The type of hydrogen fusion process that dominates in a star is determined by the temperature dependency differences between the two reactions. The proton–proton chain reaction starts at temperatures about 4×106 K,[24] making it the dominant fusion mechanism in smaller stars. A self-maintaining CNO chain requires a higher temperature of approximately 16×106 K, but thereafter it increases more rapidly in efficiency as the temperature rises, than does the proton-proton reaction.[25] Above approximately 17×106 K, the CNO cycle becomes the dominant source of energy. This temperature is achieved in the cores of main sequence stars with at least 1.3 times the mass of the Sun.[26] The Sun itself has a core temperature of about 15.7×106 K. As a main sequence star ages, the core temperature will rise, resulting in a steadily increasing contribution from its CNO cycle.[22]

Helium fusion

Main sequence stars accumulate helium in their cores as a result of hydrogen fusion, but the core does not become hot enough to initiate helium fusion. Helium fusion first begins when a star leaves the red giant branch after accumulating sufficient helium in its core to ignite it. In stars around the mass of the sun, this begins at the tip of the red giant branch with a helium flash from a degenerate helium core and the star moves to the horizontal branch where it burns helium in its core. More massive stars ignite helium in their cores without a flash and execute a blue loop before reaching the asymptotic giant branch. Despite the name, stars on a blue loop from the red giant branch are typically not blue in color, but are rather yellow giants, possibly Cepheid variables. They fuse helium until the core is largely carbon and oxygen. The most massive stars become supergiants when they leave the main sequence and quickly start helium fusion as they become red supergiants. After helium is exhausted in the core of a star, it will continue in a shell around the carbon-oxygen core.[18][21]

In all cases, helium is fused to carbon via the triple-alpha process. This can then form oxygen, neon, and heavier elements via the alpha process. In this way, the alpha process preferentially produces elements with even numbers of protons by the capture of helium nuclei. Elements with odd numbers of protons are formed by other fusion pathways.

Reaction rate

The reaction rate per volume between species A and B, having number densities nA,B is given by:

where σ(v) is the cross section at relative velocity v, and averaging is performed over all velocities.

Semi-classically, the cross section is proportional to , where is the de Broglie wavelength. Thus semi-classically the cross section is proportional to .

However, since the reaction involves quantum tunneling, there is an exponential damping at low energies that depends on Gamow factor EG, giving:

where S(E) depends on the details of the nuclear interaction.

One then integrates over all energies to get the total reaction rate, using the Maxwell–Boltzmann distribution and the relation :

where is the reduced mass.

Since this integration has an exponential damping at high energies of the form and at high energies from the Gamow factor, the integral almost vanished everywhere except around the peak, called Gamow peak, at E0, where:

Thus:

The exponent can then be approximated around E0 as:

And the reaction rate is approximated as[27]:

Values of S(E0) are typically 10-3-103 in units of keV*b, but are damped by a huge factor when involving a beta decay, due to the relation between the intermediate bound state (e.g. diproton) half-life and the beta decay half-life , as in the proton–proton chain reaction. Note that typical core temperatures in main-sequence stars give kT of the order of keV.

Thus, the limiting reaction in the CNO cycle, proton capture by 14
7
N
, has S(E0) ~ S(0) = 3.5 keV b, while the limiting reaction in the proton-proton chain reaction, the creation of deuterium from two protons, has a much lower S(E0) ~ S(0) = 4*10-22 keV b.[28][29] Incidentally, since the former reaction has a much higher Gamow factor, and due to the relative abundance of elements in typical stars, the two reaction rates are equal at a temperature value that is within the core temperature ranges of main-sequence stars.

References

  1. ^ a b Hoyle, F. (1946). "The synthesis of the elements from hydrogen". Monthly Notices of the Royal Astronomical Society. 106 (5): 343–383. Bibcode:1946MNRAS.106..343H. doi:10.1093/mnras/106.5.343.
  2. ^ a b Hoyle, F. (1954). "On Nuclear Reactions Occurring in Very Hot STARS. I. The Synthesis of Elements from Carbon to Nickel". The Astrophysical Journal Supplement Series. 1: 121. Bibcode:1954ApJS....1..121H. doi:10.1086/190005.
  3. ^ a b Burbidge, E. M.; Burbidge, G. R.; Fowler, W.A.; Hoyle, F. (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.
  4. ^ Suess, H. E.; Urey, H. C. (1956). "Abundances of the Elements". Reviews of Modern Physics. 28 (1): 53–74. Bibcode:1956RvMP...28...53S. doi:10.1103/RevModPhys.28.53.
  5. ^ a b Clayton, D. D. (1968). Principles of Stellar Evolution and Nucleosynthesis. University of Chicago Press.
  6. ^ Eddington, A. S. (1920). "The internal constitution of the stars". The Observatory. 43: 341–358. Bibcode:1920Obs....43..341E.
  7. ^ Eddington, A. S (1920). "The Internal Constitution of the Stars". Nature. 106 (2653): 14. Bibcode:1920Natur.106...14E. doi:10.1038/106014a0.
  8. ^ Selle, D. (October 2012). "Why the Stars Shine" (PDF). Guidestar. Houston Astronomical Society. pp. 6–8. Archived (PDF) from the original on 2013-12-03.
  9. ^ Bethe, H. A. (1939). "Energy Production in Stars". Physical Review. 55 (5): 434–456. Bibcode:1939PhRv...55..434B. doi:10.1103/PhysRev.55.434.
  10. ^ Lang, K. R., Life and Death of Stars, The (Cambridge: Cambridge University Press, 2013), p. 167.
  11. ^ Clayton, D. D. (2007). "HISTORY OF SCIENCE: Hoyle's Equation". Science. 318 (5858): 1876–1877. doi:10.1126/science.1151167. PMID 18096793.
  12. ^ Cameron, A. G. W. (1957). Stellar Evolution, Nuclear Astrophysics, and Nucleogenesis (PDF) (Report). Atomic Energy of Canada. Report CRL-41.
  13. ^ 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.
  14. ^ Seeger, P. A.; Fowler, W. A.; Clayton, D. D. (1965). "Nucleosynthesis of Heavy Elements by Neutron Capture". The Astrophysical Journal Supplement Series. 11: 121–126. Bibcode:1965ApJS...11..121S. doi:10.1086/190111.
  15. ^ Bodansky, D.; Clayton, D. D.; Fowler, W. A. (1968). "Nucleosynthesis During Silicon Burning". Physical Review Letters. 20 (4): 161–164. Bibcode:1968PhRvL..20..161B. doi:10.1103/PhysRevLett.20.161.
  16. ^ Bodansky, D.; Clayton, D. D.; Fowler, W. A. (1968). "Nuclear Quasi-Equilibrium during Silicon Burning". The Astrophysical Journal Supplement Series. 16: 299. Bibcode:1968ApJS...16..299B. doi:10.1086/190176.
  17. ^ Clayton, D. D. (1964). "Cosmoradiogenic Chronologies of Nucleosynthesis". The Astrophysical Journal. 139: 637. Bibcode:1964ApJ...139..637C. doi:10.1086/147791.
  18. ^ a b Jones, Lauren V. (2009), Stars and galaxies, Greenwood guides to the universe, ABC-CLIO, pp. 65–67, ISBN 978-0-313-34075-8
  19. ^ a b c d Böhm-Vitense, Erika (1992), Introduction to Stellar Astrophysics, 3, Cambridge University Press, pp. 93–100, ISBN 978-0-521-34871-3
  20. ^ Reiners, A.; Basri, G. (March 2009). "On the magnetic topology of partially and fully convective stars". Astronomy and Astrophysics. 496 (3): 787–790. arXiv:0901.1659. Bibcode:2009A&A...496..787R. doi:10.1051/0004-6361:200811450.
  21. ^ a b de Loore, Camiel W. H.; Doom, C. (1992), Structure and evolution of single and binary stars, Astrophysics and space science library, 179, Springer, pp. 200–214, ISBN 978-0-7923-1768-5
  22. ^ a b c Jeffrey, C. Simon (2010), Goswami, A.; Reddy, B. E., eds., "Principles and Perspectives in Cosmochemistry", Astrophysics and Space Science Proceedings, Springer, 16: 64–66, Bibcode:2010ASSP...16.....G, doi:10.1007/978-3-642-10352-0, ISBN 978-3-642-10368-1 |contribution= ignored (help)
  23. ^ Karttunen, Hannu; Oja, Heikki (2007), Fundamental astronomy (5th ed.), Springer, p. 247, ISBN 978-3-540-34143-7
  24. ^ Reid, I. Neill; Hawley, Suzanne L. (2005), New light on dark stars: red dwarfs, low-mass stars, brown dwarfs, Springer-Praxis books in astrophysics and astronomy (2nd ed.), Springer, p. 108, ISBN 978-3-540-25124-8
  25. ^ Salaris, Maurizio; Cassisi, Santi (2005), Evolution of stars and stellar populations, John Wiley and Sons, pp. 119–123, ISBN 978-0-470-09220-0
  26. ^ Schuler, S. C.; King, J. R.; The, L.-S. (2009), "Stellar Nucleosynthesis in the Hyades Open Cluster", The Astrophysical Journal, 701 (1): 837–849, arXiv:0906.4812, Bibcode:2009ApJ...701..837S, doi:10.1088/0004-637X/701/1/837
  27. ^ University College London astrophysics course: lecture 7 - Stars
  28. ^ Adelberger, E. G., Austin, S. M., Bahcall, J. N., Balantekin, A. B., Bogaert, G., Brown, L. S., ... & Duba, C. A. (1998). Solar fusion cross sections. Reviews of Modern Physics, 70(4), 1265.
  29. ^ Adelberger, E. G., García, A., Robertson, R. H., Snover, K. A., Balantekin, A. B., Heeger, K., ... & Chen, J. W. (2011). Solar fusion cross sections. II. The p p chain and CNO cycles. Reviews of Modern Physics, 83(1), 195.

Further reading

External links

Alastair G. W. Cameron

Alastair G. W. (Graham Walter) Cameron (21 June 1925 in Winnipeg, Manitoba, Canada – 3 October 2005 in Tucson, Arizona, USA) was a Canadian astrophysicist and space scientist who was an eminent staff member of the Astronomy department of Harvard University. He was one of the founders of the field of nuclear astrophysics, advanced the theory that the Moon was created by the giant impact of a Mars-sized object with the early Earth, and was a early adopter of computer technology in astrophysics.

Alpha nuclide

An alpha nuclide is a nuclide that consists of an integer number of alpha particles. Alpha nuclides have equal, even numbers of protons and neutrons; they are important in stellar nucleosynthesis since the energetic environment within stars is amenable to fusion of alpha particles into heavier nuclei. Alpha nuclides can be stable or radioactive. Stable alpha nuclides, and stable decay end-products of radioactive alpha nuclides, are some of the most common metals in the universe.

Alpha nuclide is also shorthand for alpha radionuclide, referring to those radioactive isotopes that undergo alpha decay and thereby emit alpha particles.

Andrew M. Davis

Andrew M. Davis is a professor of Astronomy and Geophysical Sciences at the University of Chicago. He is the son of American chemist and physicist Raymond Davis, Jr., a Nobel Prize in Physics laureate.

His main field of study is the origin of the elements by stellar nucleosynthesis. He currently is the head of a project to build a new instrument called the ion nanoprobe, which will allow isotopic and chemical analysis at finer scales than any contemporary instrument.

He is also studying the cometary dust and contemporary interstellar dust returned to Earth by the Stardust spacecraft in 2006. In 2018, he was made Fellow of the American Association for the Advancement of Science.

Argon

Argon is a chemical element with symbol Ar and atomic number 18. It is in group 18 of the periodic table and is a noble gas. Argon is the third-most abundant gas in the Earth's atmosphere, at 0.934% (9340 ppmv). It is more than twice as abundant as water vapor (which averages about 4000 ppmv, but varies greatly), 23 times as abundant as carbon dioxide (400 ppmv), and more than 500 times as abundant as neon (18 ppmv). Argon is the most abundant noble gas in Earth's crust, comprising 0.00015% of the crust.

Nearly all of the argon in the Earth's atmosphere is radiogenic argon-40, derived from the decay of potassium-40 in the Earth's crust. In the universe, argon-36 is by far the most common argon isotope, as it is the most easily produced by stellar nucleosynthesis in supernovas.

The name "argon" is derived from the Greek word ἀργόν, neuter singular form of ἀργός meaning "lazy" or "inactive", as a reference to the fact that the element undergoes almost no chemical reactions. The complete octet (eight electrons) in the outer atomic shell makes argon stable and resistant to bonding with other elements. Its triple point temperature of 83.8058 K is a defining fixed point in the International Temperature Scale of 1990.

Argon is produced industrially by the fractional distillation of liquid air. Argon is mostly used as an inert shielding gas in welding and other high-temperature industrial processes where ordinarily unreactive substances become reactive; for example, an argon atmosphere is used in graphite electric furnaces to prevent the graphite from burning. Argon is also used in incandescent, fluorescent lighting, and other gas-discharge tubes. Argon makes a distinctive blue-green gas laser. Argon is also used in fluorescent glow starters.

B2FH paper

The B2FH paper, named after the initials of the authors of the paper, Margaret Burbidge, Geoffrey Burbidge, William A. Fowler, and Fred Hoyle, is a landmark paper on the origin of the chemical elements published in Reviews of Modern Physics in 1957. The title of that paper is "Synthesis of the Elements in Stars", but as that paper grew in influence it came to be referred to only as "B2FH". The B2FH paper spread stellar nucleosynthesis theory widely in the scientific community, especially among astronomers who saw everyday relevance to their quest, at a time when it was appreciated by only a handful of experts in nuclear physics. But it did not create the theory of stellar nucleosynthesis as much as bring it vividly to life.

The paper comprehensively outlined and analyzed several key processes that are responsible for the nucleosynthesis of the elements heavier than iron and their relative abundance by the capture within stars of free neutrons. It advanced much less the understanding of the synthesis of the very abundant elements from silicon to nickel. A puzzle about that is that despite Hoyle's coauthorship of B2FH and being its chief conceptual architect, the paper did not include the carbon-burning process, the oxygen-burning process and the silicon-burning process, each of which contributes massively to the growth of stellar metallicity from magnesium to nickel in the interstellar gas. The supernova nucleosynthesis that achieves that had been published by Hoyle in 1954. Donald D. Clayton has attributed the severe undercitations of Hoyle's 1954 paper relative to the voluminous citations of B2FH to several factors: the advanced difficulty of digesting Hoyle's 1954 paper even for his B2FH coauthors, as it proved to be for the world of astronomy generally; to Hoyle's having described its key equation only in words rather than writing it prominently in his paper; and finally to a lack of careful review by Hoyle himself of the B2FH draft written by two junior coauthors who had themselves not adequately digested Hoyle's paper.

Big Bang nucleosynthesis

In physical cosmology, Big Bang nucleosynthesis (abbreviated BBN, also known as primordial nucleosynthesis, arch(a)eonucleosynthesis, archonucleosynthesis, protonucleosynthesis and pal(a)eonucleosynthesis) refers to the production of nuclei other than those of the lightest isotope of hydrogen (hydrogen-1, 1H, having a single proton as a nucleus) during the early phases of the Universe. Primordial nucleosynthesis is believed by most cosmologists to have taken place in the interval from roughly 10 seconds to 20 minutes after the Big Bang, and is calculated to be responsible for the formation of most of the universe's helium as the isotope helium-4 (4He), along with small amounts of the hydrogen isotope deuterium (2H or D), the helium isotope helium-3 (3He), and a very small amount of the lithium isotope lithium-7 (7Li). In addition to these stable nuclei, two unstable or radioactive isotopes were also produced: the heavy hydrogen isotope tritium (3H or T); and the beryllium isotope beryllium-7 (7Be); but these unstable isotopes later decayed into 3He and 7Li, as above.

Essentially all of the elements that are heavier than lithium were created much later, by stellar nucleosynthesis in evolving and exploding stars.

Chemical element

A chemical element is a species of atom having the same number of protons in their atomic nuclei (that is, the same atomic number, or Z). For example, the atomic number of oxygen is 8, so the element oxygen consists of all atoms which have exactly 8 protons.

118 elements have been identified, of which the first 94 occur naturally on Earth with the remaining 24 being synthetic elements. There are 80 elements that have at least one stable isotope and 38 that have exclusively radionuclides, which decay over time into other elements. Iron is the most abundant element (by mass) making up Earth, while oxygen is the most common element in the Earth's crust.Chemical elements constitute all of the ordinary matter of the universe. However astronomical observations suggest that ordinary observable matter makes up only about 15% of the matter in the universe: the remainder is dark matter; the composition of this is unknown, but it is not composed of chemical elements.

The two lightest elements, hydrogen and helium, were mostly formed in the Big Bang and are the most common elements in the universe. The next three elements (lithium, beryllium and boron) were formed mostly by cosmic ray spallation, and are thus rarer than heavier elements. Formation of elements with from 6 to 26 protons occurred and continues to occur in main sequence stars via stellar nucleosynthesis. The high abundance of oxygen, silicon, and iron on Earth reflects their common production in such stars. Elements with greater than 26 protons are formed by supernova nucleosynthesis in supernovae, which, when they explode, blast these elements as supernova remnants far into space, where they may become incorporated into planets when they are formed.The term "element" is used for atoms with a given number of protons (regardless of whether or not they are ionized or chemically bonded, e.g. hydrogen in water) as well as for a pure chemical substance consisting of a single element (e.g. hydrogen gas). For the second meaning, the terms "elementary substance" and "simple substance" have been suggested, but they have not gained much acceptance in English chemical literature, whereas in some other languages their equivalent is widely used (e.g. French corps simple, Russian простое вещество). A single element can form multiple substances differing in their structure; they are called allotropes of the element.

When different elements are chemically combined, with the atoms held together by chemical bonds, they form chemical compounds. Only a minority of elements are found uncombined as relatively pure minerals. Among the more common of such native elements are copper, silver, gold, carbon (as coal, graphite, or diamonds), and sulfur. All but a few of the most inert elements, such as noble gases and noble metals, are usually found on Earth in chemically combined form, as chemical compounds. While about 32 of the chemical elements occur on Earth in native uncombined forms, most of these occur as mixtures. For example, atmospheric air is primarily a mixture of nitrogen, oxygen, and argon, and native solid elements occur in alloys, such as that of iron and nickel.

The history of the discovery and use of the elements began with primitive human societies that found native elements like carbon, sulfur, copper and gold. Later civilizations extracted elemental copper, tin, lead and iron from their ores by smelting, using charcoal. Alchemists and chemists subsequently identified many more; all of the naturally occurring elements were known by 1950.

The properties of the chemical elements are summarized in the periodic table, which organizes the elements by increasing atomic number into rows ("periods") in which the columns ("groups") share recurring ("periodic") physical and chemical properties. Save for unstable radioactive elements with short half-lives, all of the elements are available industrially, most of them in low degrees of impurities.

Chemical evolution

Chemical evolution may refer to:

Abiogenesis, the transition from nonliving elements to living systems

Astrochemistry, the study of the abundance and reactions of molecules in the Universe, and their interaction with radiation

Cosmochemistry, the study of the chemical compositions in the universe and the processes that led to them

Evolution of metal ions in biological systems, incorporation of metallic ions into living organisms and how it has changed over time

Gas evolution reaction, the process of a gas bubbling out from a solution

Molecular evolution, evolution at the scale of molecules

Oxygen evolution, the process of generating molecular oxygen through chemical reaction

Stellar nucleosynthesis, the creation of chemical elements by stellar thermonuclear fusion or supernovae

Fred Hoyle

Sir Fred Hoyle FRS (24 June 1915 – 20 August 2001) was a British astronomer who formulated the theory of stellar nucleosynthesis. He also held controversial stances on other scientific matters—in particular his rejection of the "Big Bang" theory, a term coined by him on BBC radio, and his promotion of panspermia as the origin of life on Earth. He also wrote science fiction novels, short stories and radio plays, and co-authored twelve books with his son, Geoffrey Hoyle.

He spent most of his working life at the Institute of Astronomy at Cambridge and served as its director for six years.

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.

Iron peak

The iron peak is a local maximum in the vicinity of Fe (Cr, Mn, Fe, Co and Ni) on the graph of the abundances of the chemical elements, as seen below.

For elements lighter than iron on the periodic table, nuclear fusion releases energy while fission consumes it. For iron, and for all of the heavier elements, nuclear fusion consumes energy, but nuclear fission releases it. Chemical elements up to the iron peak are produced in ordinary stellar nucleosynthesis. Heavier elements are produced only during supernova nucleosynthesis. This is why we have more iron peak elements than in its neighbourhood.

Metallicity

In astronomy, metallicity is used to describe the abundance of elements present in an object that are heavier than hydrogen or helium. Most of the physical matter in the Universe is in the form of hydrogen and helium, so astronomers use the word "metals" as a convenient short term for "all elements except hydrogen and helium". This usage is distinct from the usual physical definition of a solid metal. For example, stars and nebulae with relatively high abundances of carbon, nitrogen, oxygen, and neon are called "metal-rich" in astrophysical terms, even though those elements are non-metals in chemistry.

The presence of heavier elements hails from stellar nucleosynthesis, the theory that the majority of elements heavier than hydrogen and helium in the Universe ("metals", hereafter) are formed in the cores of stars as they evolve. Over time, stellar winds and supernovae deposit the metals into the surrounding environment, enriching the interstellar medium and providing recycling materials for the birth of new stars. It follows that older generations of stars, which formed in the metal-poor early Universe, generally have lower metallicities than those of younger generations, which formed in a more metal-rich Universe.

Observed changes in the chemical abundances of different types of stars, based on the spectral peculiarities that were later attributed to metallicity, led astronomer Walter Baade in 1944 to propose the existence of two different populations of stars.

These became commonly known as Population I (metal-rich) and Population II (metal-poor) stars. A third stellar population was introduced in 1978, known as Population III stars. These extremely metal-poor stars were theorised to have been the "first-born" stars created in the Universe.

Neutron capture nucleosynthesis

Neutron capture nucleosynthesis describes two nucleosynthesis pathways: the r-process and the s-process, for rapid and slow neutron captures, respectively. R-process describes neutron capture in a region of high neutron flux, such as during supernova nucleosynthesis after core-collapse, and yields neutron-rich nuclides. S-process describes neutron capture that is slow relative to the rate of beta decay, as for stellar nucleosynthesis in some stars, and yields nuclei with stable nuclear shells. Each process is responsible for roughly half of the observed abundances of elements heavier than iron. The importance of neutron capture to the observed abundance of the chemical elements was first described in 1957 in the B2FH paper.

Nuclear transmutation

Nuclear transmutation is the conversion of one chemical element or an isotope into another chemical element. Because any element (or isotope of one) is defined by its number of protons (and neutrons) in its atoms, i.e. in the atomic nucleus, nuclear transmutation occurs in any process where the number of protons or neutrons in the nucleus is changed.

A transmutation can be achieved either by nuclear reactions (in which an outside particle reacts with a nucleus) or by radioactive decay, where no outside cause is needed.

Natural transmutation by stellar nucleosynthesis in the past created most of the heavier chemical elements in the known existing universe, and continues to take place to this day, creating the vast majority of the most common elements in the universe, including helium, oxygen and carbon. Most stars carry out transmutation through fusion reactions involving hydrogen and helium, while much larger stars are also capable of fusing heavier elements up to iron late in their evolution.

Elements heavier than iron, such as gold and lead, are created through elemental transmutations that can only take place in supernovae - as stars begin to fuse heavier elements, substantially less energy is released from each fusion reaction, and each fusion reaction that produces elements heavier than iron is endothermic in nature, and stars are incapable of carrying this out.

One type of natural transmutation observable in the present occurs when certain radioactive elements present in nature spontaneously decay by a process that causes transmutation, such as alpha or beta decay. An example is the natural decay of potassium-40 to argon-40, which forms most of the argon in the air. Also on Earth, natural transmutations from the different mechanisms of natural nuclear reactions occur, due to cosmic ray bombardment of elements (for example, to form carbon-14), and also occasionally from natural neutron bombardment (for example, see natural nuclear fission reactor).

Artificial transmutation may occur in machinery that has enough energy to cause changes in the nuclear structure of the elements. Such machines include particle accelerators and tokamak reactors. Conventional fission power reactors also cause artificial transmutation, not from the power of the machine, but by exposing elements to neutrons produced by fission from an artificially produced nuclear chain reaction. For instance, when a uranium atom is bombarded with slow neutrons, fission takes place. This releases, on average, 3 neutrons and a large amount of energy. The released neutrons then cause fission of other uranium atoms, until all of the available uranium is exhausted. This is called a chain reaction.

Artificial nuclear transmutation has been considered as a possible mechanism for reducing the volume and hazard of radioactive waste.

Nucleosynthesis

Nucleosynthesis is the process that creates new atomic nuclei from pre-existing nucleons, primarily protons and neutrons. The first nuclei were formed about three minutes after the Big Bang, through the process called Big Bang nucleosynthesis. Seventeen minutes later the universe had cooled to a point at which these processes ended, so only the fastest and simplest reactions occurred, leaving our universe containing about 75% hydrogen, 24% helium, and traces of other elements such as lithium and the hydrogen isotope deuterium. The universe still has approximately the same composition today.

Heavier nuclei were created from these, by several processes. Stars formed, and began to fuse light elements to heavier ones in their cores, giving off energy in the process, known as stellar nucleosynthesis. Fusion processes create many of the lighter elements up to and including iron and nickel, and these elements are ejected into space (the interstellar medium) when smaller stars shed their outer envelopes and become smaller stars known as white dwarfs. The remains of their ejected mass form the planetary nebulae observable throughout our galaxy.

Supernova nucleosynthesis within exploding stars by fusing carbon and oxygen is responsible for the abundances of elements between magnesium (atomic number 12) and nickel (atomic number 28). Supernova nucleosynthesis is also thought to be responsible for the creation of rarer elements heavier than iron and nickel, in the last few seconds of a type II supernova event. The synthesis of these heavier elements absorbs energy (endothermic process) as they are created, from the energy produced during the supernova explosion. Some of those elements are created from the absorption of multiple neutrons (the r-process) in the period of a few seconds during the explosion. The elements formed in supernovas include the heaviest elements known, such as the long-lived elements uranium and thorium.

Neutron star mergers and collisions are also responsible for many heavy elements, via the r-process ("r" stands for "rapid"). Neutron stars are extremely dense remnants of supernovae, and as their name suggests, they consist of a complex state of matter, predominantly made of tightly packed neutrons. When two such dense stars collide, a large amount of neutron-rich matter may be ejected at extremely high temperatures and under exotic conditions, and heavy elements may form as the ejecta begins to cool. In 2017, the merger of GW170817 led to the detection of substantial signatures of gold, platinum and other heavy elements over an extended period.

Cosmic ray spallation, caused when cosmic rays impact the interstellar medium and fragment larger atomic species, is a significant source of the lighter nuclei, particularly 3He, 9Be and 10,11B, that are not created by stellar nucleosynthesis.

In addition to the fusion processes responsible for the growing abundances of elements in the universe, a few minor natural processes continue to produce very small numbers of new nuclides on Earth. These nuclides contribute little to their abundances, but may account for the presence of specific new nuclei. These nuclides are produced via radiogenesis (decay) of long-lived, heavy, primordial radionuclides such as uranium and thorium. Cosmic ray bombardment of elements on Earth also contribute to the presence of rare, short-lived atomic species called cosmogenic nuclides.

Presolar grains

Presolar grains are interstellar solid matter in the form of tiny solid grains that originated at a time before the Sun was formed (presolar: before the Sun). Meteoriticists often use the term to represent stardust, grains that originated within a single star and which they extract from meteorites for study. Because most interstellar grains are not stardust from a single star, however, being instead interstellar cloud matter accreted by smaller presolar grains, most presolar grains are also not stardust. Logically, all stardust are presolar grains; but not all presolar grains are stardust. This confusing terminology is heavily entrenched among 21st century meteoriticists who prefer to use the terms interchangeably, however, so it is best to live with both usages or to write presolar stardust grains for stardust.

Presolar stardust grains formed within outflowing and cooling gases from earlier presolar stars. The stellar nucleosynthesis that took place within each presolar star gives to each granule an isotopic composition unique to that parent star, which differs from the isotopic composition of our solar system's matter as well as from the galactic average. These isotopic signatures often fingerprint very specific astrophysical nuclear processes that took place within the parent star and prove their presolar origin.

Rubidium–strontium dating

The rubidium-strontium dating method is a radiometric dating technique used by scientists to determine the age of rocks and minerals from the quantities they contain of specific isotopes of rubidium (87Rb) and strontium (87Sr, 86Sr).

Development of this process was aided by German chemists Otto Hahn and Fritz Strassmann, who later went on to discover nuclear fission in December 1938.

The utility of the rubidium–strontium isotope system results from the fact that 87Rb (one of two naturally occurring isotopes of rubidium) decays to 87Sr with a half-life of 49.23 billion years. In addition, Rb is a highly incompatible element that, during partial melting of the mantle, prefers to join the magmatic melt rather than remain in mantle minerals. As a result, Rb is enriched in crustal rocks. The radiogenic daughter, 87Sr, is produced in this decay process and was produced in rounds of stellar nucleosynthesis predating the creation of the Solar System.

Different minerals in a given geologic setting can acquire distinctly different ratios of radiogenic strontium-87 to naturally occurring strontium-86 (87Sr/86Sr) through time; and their age can be calculated by measuring the 87Sr/86Sr in a mass spectrometer, knowing the amount of 87Sr present when the rock or mineral formed, and calculating the amount of 87Rb from a measurement of the Rb present and knowledge of the 85Rb/87Rb weight ratio.

If these minerals crystallized from the same silicic melt, each mineral had the same initial 87Sr/86Sr as the parent melt. However, because Rb substitutes for K in minerals and these minerals have different K/Ca ratios, the minerals will have had different Rb/Sr ratios.

During fractional crystallization, Sr tends to become concentrated in plagioclase, leaving Rb in the liquid phase. Hence, the Rb/Sr ratio in residual magma may increase over time, resulting in rocks with increasing Rb/Sr ratios with increasing differentiation.

Highest ratios (10 or higher) occur in pegmatites.

Typically, Rb/Sr increases in the order plagioclase, hornblende, K-feldspar, biotite, muscovite. Therefore, given sufficient time for significant production (ingrowth) of radiogenic 87Sr, measured 87Sr/86Sr values will be different in the minerals, increasing in the same order.

Supernova nucleosynthesis

Supernova nucleosynthesis is a theory of the nucleosynthesis of the natural abundances of the chemical elements in supernova explosions, advanced as the nucleosynthesis of elements from carbon to nickel in massive stars by Fred Hoyle in 1954. In massive stars, the nucleosynthesis by fusion of lighter elements into heavier ones occurs during sequential hydrostatic burning processes called helium burning, carbon burning, oxygen burning, and silicon burning, in which the ashes of one nuclear fuel become, after compressional heating, the fuel for the subsequent burning stage. During hydrostatic burning these fuels synthesize overwhelmingly the alpha-nucleus (A = 2Z) products. A rapid final explosive burning is caused by the sudden temperature spike owing to passage of the radially moving shock wave that was launched by the gravitational collapse of the core. W. D. Arnett and his Rice University colleagues demonstrated that the final shock burning would synthesize the non-alpha-nucleus isotopes more effectively than hydrostatic burning was able to do, suggesting that the expected shock-wave nucleosynthesis is an essential component of supernova nucleosynthesis. Together, shock-wave nucleosynthesis and hydrostatic-burning processes create most of the isotopes of the elements carbon (Z = 6), oxygen (Z = 8), and elements with Z = 10–28 (from neon to nickel). As a result of the ejection of the newly synthesized isotopes of the chemical elements by supernova explosions their abundances steadily increased within interstellar gas. That increase became evident to astronomers from the initial abundances in newly born stars exceeding those in earlier-born stars. To explain that temporal increase of the natural abundances of the elements was the main goal of stellar nucleosynthesis. Hoyle's paper was the founding paper of that theory; however, ideas about nuclear reactions in stars providing power for the stars is often confused with stellar nucleosynthesis. Realize that nuclear fusion in stars can occur with negligible impact on the abundances of the chemical elements.

Elements heavier than nickel are comparatively rare owing to the decline with atomic weight of their nuclear binding energies per nucleon, but they too are created in part within supernovae. Of greatest interest historically has been their synthesis by rapid capture of neutrons during the r-process, reflecting the common belief that supernova cores are likely to provide the necessary conditions. But see the r-process below for a recently discovered alternative. The r-process isotopes are roughly a 100,000 times less abundant than the primary chemical elements fused in supernova shells above. Furthermore, other nucleosynthesis processes in supernovae are thought to also be responsible for some nucleosynthesis of other heavy elements, notably, the proton capture process known as the rp-process, the slow capture of neutrons (s-process) in the Helium-burning shells and in the carbon-burning shells of massive stars, and a photodisintegration process known as the γ-process (gamma-process). The latter synthesizes the lightest, most neutron-poor, isotopes of the elements heavier than iron from preexisting heavier isotopes.

Triple-alpha process

The triple-alpha process is a set of nuclear fusion reactions by which three helium-4 nuclei (alpha particles) are transformed into carbon.

This page is based on a Wikipedia article written by authors (here).
Text is available under the CC BY-SA 3.0 license; additional terms may apply.
Images, videos and audio are available under their respective licenses.