Stellar nucleosynthesis

Stellar nucleosynthesis refers to the assembly of the natural abundances of the chemical elements by nuclear reactions occurring in the cores of stars. Those stars evolve (age) owing to the associated changes in the abundances of the elements within. Those stars lose most of their mass when it is ejected late in the stellar lifetimes, thereby enriching the interstellar gas in the abundances of elements heavier than helium. For the creation of elements during the explosion of a star, the term supernova nucleosynthesis is used. The goal is to understand the vastly differing abundances of the chemical elements and their several isotopes as a process of natural history. The primary stimulus to the development of this theory was the shape of the natural abundances. Those abundances, when plotted on a graph as a function of atomic number of the element, have a jagged sawtooth structure varying by factors of ten million. This suggested a natural process rather than a random distribution. Such a graph of the abundances can be seen at History of nucleosynthesis theory. Stellar nucleosynthesis is the most dominating contributor to several processes that also occur under the collective term nucleosynthesis.

A second major stimulus to understanding the processes involved occurred throughout the 20th century, when it was first realized that the energy released from nuclear fusion reactions accounted for the longevity of the Sun as a source of heat and light. The fusion of heavier nuclei from initial hydrogen and helium provides that energy source, which synthesizes new nuclei as a byproduct of the fusion. This became clear during the decade prior to WWII. Those associated fusion product nuclei are restricted to nuclei only slightly heavier than the fusing nuclei, however, and thus do not contribute heavily to the natural abundances of the elements. Nonetheless, this success raised the plausibility of explaining all of the natural abundances in this way. The prime energy producer in the sun is the fusion of hydrogen to helium, which occurs at a minimum temperature of 3 million kelvin.

 

 

History

In 1920, Arthur Eddington, on the basis of the precise measurements of atoms by F.W. Aston, was the first to suggest that stars obtained their energy from nuclear fusion of hydrogen to form helium. 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. He selected two processes that he believed to be the sources of energy in stars. The first one, the proton-proton chain, 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 most important in more massive stars. These works concerned the energy generation capable of keeping stars hot. They 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 into iron. Hoyle followed that in 1954 with a large paper describing how advanced fusion stages within stars would synthesize elements between carbon and iron in mass. This is the dominant work in stellar nucleosynthesis. It provided the roadmap to how the most abundant elements on earth had been synthesized from initial hydrogen and helium, making clear how those abundant elements increased their galactic abundances as the galaxy aged.

Quickly, many important omissions in Hoyle's theory were corrected, beginning with the publication of a celebrated review paper in 1957 by Burbidge, Burbidge, Fowler and Hoyle (commonly referred to as the B²FH paper). 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 work as much as many assumed, except in the understanding of nucleosynthesis of those elements heavier than iron. Significant improvements were made by Alastair GW Cameron and by Donald D. Clayton. Cameron presented his own independent approach (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 and of the R-process, of the burning of silicon into iron-group elements, and discovered radiogenic chronologies for determining the age of the elements. The entire research field expanded rapidly in the 1970s.