Origin of the Elements

1. Origin of the Elements Cosmochemistry I Lecture 39

2. Cosmochemistry • Questions: • When and how did the elements form? • What is the universe composed of? Is it uniform in space and time? • When and how did the solar system form? • What is the solar system composed of? • When and how did the Earth form? • Sources of answers • Laws of physics and chemistry • Astronomical observations, including spectral (chemical) observations of stars, including the Sun, and planetary objects • Meteorites • Lunar Samples • Remote analyses from Martian landers • Interplanetary dust and solar wind

3. Astronomical Background • Stars are classified based on their color (and spectral absorption lines), which is in turn related to their surface temperature. • On a plot of luminosity versus wavelength of their principal emissions (color), called, most stars fall along an array defining the “main sequence”. • Since wavelength is inversely related to the fourth power of temperature, this correlation means that hot stars give off more energy than cooler stars. Mass is also related to temperature for main sequence stars: hot stars are big, cool stars are small. • The relationship between mass, luminosity, and temperature is nonlinear. For example, an O star that is 30 times as massive as the Sun will have a surface temperature 7 times as hot and a luminosity 100,000 times brighter. • Stars on the main sequence produce energy by “hydrogen burning”, fusion of hydrogen to produce helium. The relationship results from the rate of hydrogen burning: large stars have hot, dense interiors and burn hydrogen much faster than smaller stars. • Consequently there is an inverse relationship between the main sequence lifetime of a star and mass. The most massive stars, up to ~100 solar masses, have life expectancies of only about 106 years, whereas small stars, ~0.01 solar masses, remain on the main sequence for more than 1010 years. • Stars are also divided into Population I stars, greater heavy element contents, and Population II stars, small, metal-poor, occurring mainly in globular clusters outside main galactic disk. Hertzsprung-Russell diagram

4. In the beginning… • The universe began infinitely hot and infinitesimally small 13.8 billion years ago in the Big Bang. • Since then the universe has been expanding, cooling, and evolving. • Stars were born and stars died. • Galaxies formed and merged.

5. Non-Main Sequence Stars • The two most important exceptions to the main sequence stars, the red giants and the white dwarfs, represent stars that have burned all the H fuel in the cores and have moved on in the evolutionary sequence. • H in the core is not replenished because the density difference prevents convection between the core and outer layers, which are still H-rich. The interior part of the core collapses under gravity until temperature and pressure are great enough for He burning to begin. At the same time the exterior expands and cools, resulting in a red giant, a star that is overluminous relative to main sequence stars of the same color. • The fate of stars when the He in the core is exhausted depends on their mass. Nuclear reactions in small stars cease and they simply contract, their exteriors heating up as they do so, to become white dwarfs. Larger stars go onto to ‘burn’ heavier elements in their cores – we’ll talk about that later. • The other exceptions are the T-Tauri types stars. These are young stars in which fusion has not yet ignited. They are radiating the energy of gravitational collapse.

6. Polygenetic hypothesis and the origin of the elements • Our understanding of nucleosynthesis comes from three sets of observations: • (1) the abundance of isotopes and elements in the cosmos; • (2) experiments on nuclear reactions that determine what reactions are possible (or probable) under given conditions; and • (3) inferences about possible sites of nucleosynthesis and about the conditions that prevail in those sites. • Two possibilities for formation of the elements: • (1) they were formed in the Big Bang itself • (2) they were subsequently produced • Differences between old Population II stars and younger Population I stars suggests the universe has evolved. • However, no single mechanism could be found to explain the observed abundances of the elements.

7. B2FH and Nucleosynthesis • Burbidge, Burbidge, Fowler and Hoyle (1957) proposed the elements were created in 4 ways/environments: • Cosmological nucleosynthesis: creation in the Big Bang • Stellar nucleosynthesis: synthesis of elements by fusion in stars • Explosive nucleosynthesis: synthesis of elements by neutron and proton capture reactions in supernovae • Galactic nucleosynthesis: synthesis of elements by cosmic ray spallation reactions Margaret and Geoffrey Burbidge

8. Cosmological Nucleosynthesis • Immediately after the Big Bang, the universe was too hot for any matter to exist. • But within a microsecond or so, it had cooled to 1011 K so that matter began to condense. • At first electrons, positrons, and neutrinos dominated, but as the universe cooled and expanded, protons and neutrons became more abundant. These existed in an equilibrium dictated by the following reactions: 1H + e– ⇄ n + ν n + e+ ⇄ 1H + ν • As temperatures cooled through 1010 K, the reactions above progressively favored protons. In less than two seconds things had cooled enough so that these reactions ceased, freezing in a 6 to 1 ratio of protons to neutrons. • It took another 100 seconds for the universe to cool to 109 K, which is cool enough for 2H to form: 1H + 1n ⇋ 2H + γ • Subsequent reactions produced 3He, 4He and a wee bit of Li. • Within 20 minutes or so, the universe cooled below 3 x108 K and nuclear reactions were no longer possible. • Some 400,000 years later, the universe had cooled to about 3000 K, cool enough for electrons to be bound to nuclei, forming atoms.

9. Stellar Nucleosynthesis • When density of a forming star reaches 6 g/cm and T reached 10 to 20 million K, hydrogen burning, or the pp process, can begin which involves reactions such as: 1H + 1H → 2H + β++ ν 2H + 1H → 3He + γ 3He + 3He → 4He + 21H + γ • CNO cycle: carbon acts a nuclear catalyst to also synthesize 4He from 1H • 12C(p,γ) 13N(β++,γ) 13C(p,γ) 14N(p,γ) 15O(β+,ν) 15N(p,α) 12C • limited to larger Pop. I stars • These are the sources of energy sustaining main sequence stars. • Little synthesis beyond He; some minor production/consumption of light nuclides, particularly in the CNO cycle.

10. Clarification: AFC equation • Equation written as: • is for Rdefined as the ratio of mass assimilated to mass crystallized (as shown in Figure). • For R defined as mass crystallized to mass assimilated, the equation should be • (but this is not how the figure is labeled).

11. Example: p, r, and s nuclides

12. Stellar Nucleosynthesis in Red Giants • Once the H is exhausted in the stellar core the interior collapses, raising T and P. • The exterior expands and cools. This is the red giant phase. • When T reaches 108 K and density reaches 104 g/cc in the He core), He burning begins: 4He + 4He → 8Be + γ 8Be + 4He → 12C + γ • Because the t1/2of 8Be is only 10-16sec, 3He must collide effectively simultaneously, which is why pressure must be so high. • He burning also produces some O, 20Ne and 24Mg but Li, Be, and B are skipped: they are not synthesized, rather they are consumed in stars. • Once He is consumed in the core, low mass stars such as the Sun cannot reach T and P for heavier fusion reactions and they end their lives as white dwarfs. • Stars bigger than about 4 M☼ undergo further collapse and the initiation of carbon burningwhen temperatures reach 600 million K and densities 5 x105 g/cc. • For stars more massive than 11 M☼, about 1% of all stars, evolution now proceeds at an exponentially increasing pace as successive fusion reactions at higher T and P. Evolution of a 25 solar mass star.

13. The e-process • As the finale approaches, the star has become a cosmic onion of sorts, with layers of heavier and heavier elements. • A new core consisting mainly of 28Si has been created. • At temperatures near 109 K and densities above 107 g/cc a process known as silicon burning, or the e-process (for equilibrium). • This process is really a variety of reactions that can be summarized as the photonuclear rearrangement of a gas originally consisting of 28Si nuclei into one which consists mainly of 56Ni, which then decays with a half-life of 6 days to 56Fe, the most stable of all nuclei. • The e-process includes reactions such as: 28Si + γ ⇄ 24Ne + 4He 28Si + 4He ⇄ 32S + γ 32S + 4He ⇄ 36Ar + γ • While these reactions can proceed in either direction, there is a tendency for the build-up of heavier nuclei with masses 32, 36, 40, 44, 48, 52, and 56, Partly as a result of the e-process, these nuclei are unusually abundant in nature. A variety of minor nuclei are produced as well. • This continues for a few days at most. Finally, the inner core has been converted completely to 56Ni and 56Fe, the latter the most stable of all nuclei. Exogenic fusion reactions are no longer possible.

14. In the meantime… The s-process • In second and later generation stars containing heavy elements, yet another nucleosynthetic process can operate. This is the slow neutron capture or s-process. It operates mainly in the red giant phase, during which some fusion reaction produce free neutrons, e.g.: • 13C + 4He → 16O + n • These neutrons are captured by nuclei to produce successively heavier elements. • The rate of production and the rate of capture, however, is slow. Consequently, a radioactive isotope of an element will decay before it can capture a second neutron and gaps between isotopes of an element cannot be bridged. This is the main difference between this and the r-process.

15. Explosive Nucleosynthesis • Once the stellar core has been largely converted to Fe, a critical phase is reached: the balance between thermal expansion and gravitational collapse is broken. The stage is now set for the catastrophic death of the star: a supernova explosion, the ultimate fate of stars with masses greater than about 8 solar masses. • The energy released in the supernova is astounding. In its first 10 seconds, the 1987A supernova released more energy than the entire visible universe, and 100 times more energy than the Sun will release in its entire 10 billion year lifetime. • The supernova begins with the collapse of the stellar core, which would have a radius similar to the Earth’s radius before collapse, collapsing to a radius of 100 km or so in a few tenths of a second. When matter in the center of the core is compressed beyond the density of nuclear matter (3 ´ 1014 g/cc), it rebounds, sending a massive shock wave back out. As the shock wave travels outward through the core, the temperature increase resulting from the compression produces a breakdown of nuclei by photodisintegration, for example: • 56Fe + γ → 13 4He + 4 1n; • 4He + γ2 1H + 2 1n • This photodisintegration produces a large number of free neutrons (and protons), which leads to another important nucleosynthetic process, the r-process.

16. The r- & p-processes • Because the abundance of neutrons is exceedingly high, nuclei capture them at a rapid rate – so rapid that even an unstable nucleus will capture a neutron before it has an opportunity to decay. The result is a build-up of neutron-rich unstable nuclei. Eventually the nuclei capture enough neutrons that they are not stable even for a small fraction of a second. At that point, they undergoβ-decay to new nuclides, which are more stable and capable of capturing more neutrons. • This is the principal mechanism for building up the heavier nuclei. • It reaches a limit when nuclei beyond A ≈ 90 are reached. These heavy nuclei fission into several lighter fragments. • The r-process is thought to have a duration of 1 to 100 sec during the peak of the supernova explosion. • The r-process tends to form the heavier (neutron-rich) isotopes of a given element. Proton capture, or the p-process, also occurs in supernovae. The probability of proton capture is much less than that of neutron capture, so it is significant only for the production of the lightest isotopes of a given element.

17. r- s- and -processes summary • Shielding: If an isotope with z protons and n neutrons has a stable isobar with n + x neutrons and p -x protons, this isotope is shielded from production by the r-process because b-decay will cease when that stable isobar is reached. • S-process path is zig-zag: there is a zag every time there is a non-stable isotope of an element: in that case, the path zags to the next higher atomic number element. • The least abundant are those created by only one, particularly by only the p-process. These tend to be both r-process shielded and off the s-process path to low N (to the left). • The most abundant isotopes of an element tend to be those created by all processes • The exact abundance of an isotope depends on a number of factors, including its neutron-capture cross-section and the neutron-capture cross-section and stability of neighboring nuclei. The neutron-capture cross-section of a nuclide is a measure of the affinity of that nuclide for neutrons, i.e., a measure of the probability of that nuclide capturing a neutron in a given neutron flux.

18. Galactic Nucleosynthesis • Except for production of 7Li in the Big Bang, Li, Be, and B are not produced in any of the above situations. • One clue to the creation of these elements is their abundance in galactic cosmic rays: they are overabundant by a factor of 106. • They are believed to be formed by interactions of cosmic rays with interstellar gas and dust, primarily reactions of 1H and 4He with carbon, nitrogen, and oxygen nuclei. • These reactions occur at high energies (higher than the Big Bang and stellar interiors), but at low temperatures where the Li, B and Be can survive.