Stellar Spectroscopy and Elemental Abundances Definitions Solar Abundances Relative Abundances Origin of Elements 1

Definitions X, Y, Z: amounts of H, He, and rest (metals) by mass (total of 1; log Z/Z sun in Kurucz models) Solar: X=0735, Y=0.248, Z=0.017 Abundance as number density relative to H A=N(element)/N(H) Usually given as log A or log A +12 [Fe/H]=log[N(Fe)/N(H)] star – log[N(Fe)/N(H)] sun [M/H] sometimes reported as mean metallicity 2

Solar Abundances From spectroscopy and meteorites Gray Table 16.3 Scott, Grevesse et al. 2014arXiv S 2014arXiv S 3

Relative Abundances 4

Origin of the Elements Hydrogen is most abundant element, followed fairly closely by helium. He formed in the Big Bang, with some increase from the primordial He abundance (Y p =0.24) by subsequent H-burning in stars (Y =0.28 here and now). 5

Light Elements: Li, Be, B Li, Be and B are very scarce, mostly destroyed in the harsh environment of stellar interiors Li abundance comes from measurements in meteorites; it is still lower in the solar photosphere because of destruction by mixing with hotter layers below. Abundant in primary cosmic rays as a result of fusion and spallation reactions between p and (mainly) CNO nuclei at high energies. Deuterium and some Li formed in Big Bang. 6

Carbon (6) to Calcium (20) Downward progression modulated by odd:even and shell effects in nuclei which affect their binding energy. From successive stages in stellar evolution: exhaustion of one fuel is followed by contraction, heating, alpha=He capture fusion. Onset of Ca burning leads to Mg and nearby elements; accompanied by neutrino emission (ever faster evolution). 7

Iron Group Fe-group elements represent approximate nuclear statistical equilibrium at T ≈10 9 K Result of shock that emerges from the core of a massive star that has collapsed into a neutron star (SN II) OR sudden ignition of C in a white dwarf that has accreted enough material from a companion to bring it over the Chandrasekhar mass limit (SN Ia). Dominant product is 56 Ni, most stable nucleus with equal numbers of protons and neutrons, which later decays into 56 Fe. 8

s-process: slow addition of neutrons Nucleosynthesis beyond the Fe group occurs neutron capture. Captures on a seed nucleus (mostly 56 Fe) lead to the production of a β- unstable nucleus (e.g. 59 Fe). Outcome depends on relative time-scales for neutron addition and decay. s-process: slow addition, so that unstable nuclei have time to undergo decay Nuclei form along the stability valley to 209 Bi. 9

s- and r-process decays in neighborhood of Tin (Sn) 10 β decay

r-process: rapid addition of neutrons Many neutrons are added under conditions of very high T, neutron density; build unstable nuclei up to the point where (n, γ) captures are balanced by (γ, n) photodisintegrations After neutron supply is switched off, products undergo a further decays ending at the nearest stable isobar (neutron-rich side of the stability valley). Some elements from both r- and s-processes. 11

Stability Valley 12 Abundance peaks occur corresponding to closed shells with 50, 82 or 126 neutrons

Summary of Relative Abundances 13

Metal rich vs. metal poor stars: Frebel et al. 2005, Nature, 434, 871 [Fe/H]=

Abundance Trends Metals higher in Pop I stars (younger, disk) than in Pop II stars (older, halo); Galactic enrichment with time Metals higher closer to Galactic center Evolutionary changes: Li decrease with age CNO-processed gas in stars with mixing C enhancement in older stars with He-burning Magnetic fields can create patches with unusual abundance patterns: Ap stars 15