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Lecture 2: Formation of the chemical elements

Lecture 2: Formation of the chemical elements. Bengt Gustafsson: Current problems in Astrophysics Ångström Laboratory, Spring 2010. Outline. Basic physics Observation methods Big Bang Nucleosynthesis, just a few words Stellar Nucleosynthesis

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Lecture 2: Formation of the chemical elements

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  1. Lecture 2: Formation of the chemical elements Bengt Gustafsson: Current problems in Astrophysics Ångström Laboratory, Spring 2010

  2. Outline • Basic physics • Observation methods • Big Bang Nucleosynthesis, just a few words • Stellar Nucleosynthesis • Inter-stellar Nucleosynthesis, just a few words

  3. Basic Physics: thermo-nuclear reactions • At LTE: Maxwell distribution • Coulomb barrier: • Cross section: or tunneling probablity = Sommerfeld parameter. S(E) is slowlyvaryying (if no resonances)

  4. Gamow peak: Note extreme T dependence!: E.g. x2 in T => x106 in yield G. Gamow (1904-68)

  5. Cross sections, resonances • Note, reactions in stars occur at relatively low energies (~ 100 keV) • Resonances still a major problem for many important reactions!

  6. Reaction networks -- stellar models

  7. Observations Solar (system) abundances

  8. Big bang nucleosynthesis • (1 s) - 3 min - 20 min • 1010 K -- 2 108 K

  9. Chemical elements • were there from the beginning? • formed in the Big Bang (Alpher, Bethe, Gamow 1948) • formed in stars (Fred Hoyle 1946) • Burbidge, Burbidge, Fowler & Hoyle (1957), ”B2FH” William Fowler 1911-1995 Fred Hoyle 1915-2001

  10. Discoveries in 1950:ies of key significance: Tc in Mira stars (Merrill 1952) -- half life 4.2 Myear Subdwarfs (Pop II) very metal-poor (Chamberlain & Aller 1951) Ba II stars (rich in s elements) Since then stellar spectroscopy developed: Spectra in high resolution, high S/N Model atmospheres to model the spectra Abundances to better than 10%-30% today.

  11. Stellar nucleosynthesis pp chains The slow part T > 107 K

  12. CNO cycle Hans Bethe 1906-2005 T > 13 106 K Transforms 12C and 16O to 14N and 13C in addition to 41H 4He + energy

  13. Edwin Salpeter 1924-2008 Trippel  process T ~ 108 K 4He + 4He → 8Be (−92 keV) 8Be + 4He → 12C + e+ + e⁻ (+7.367 MeV) Resonance required here: Fred Hoyle’s prediction!

  14. For mass > 10 Msun, carbon burning starts at 600 MK Si burning at 2.7 GK, lasts ~5 days; ”explosive” Also O burning Then, no energy left => Rapid core contraction, Photodisintegration of nuclei => Neutron star

  15. The s-process Adding neutrons to heavy nuclei: ~105-1011 n per cm2 and s In red giants (Miras, S stars, C stars, maybe Ba II stars) Alistair Cameron 1925-2005

  16. r process n fluxes of typically ~1022 per cm2 and s In principle available in supernovae

  17. SN 1987a in LMC of Type IIp How do they explode? Collapse - Bounce - - Shock wave - - Neutrinos from p+n (seen) But models do not explode … Supernovae Type II

  18. … until seemingly very recently • H. Th. Janka et al. MPI München 0.4 s, 11.2 Msun 0.7 s, 15 Msun Complex 3D instabilities!

  19. Supernovae type Ia http://www.ci.uchicago.edu/flashviz/gallery/main.php?g2_itemId=4827 • HD 3D Simulations by R. Fischer et al. • A deflagrating white dwarf (A SN type Ia) • A few seconds event • => Fe, Ni, Si, Ca • Original mass ~ 3Msun => come after Type II.

  20. Bensby & Feltzing (2008): Galactic disk stars Major Fe source (SN Ia) came after major O source (SN II). Note two different populations with different age and different kinematics

  21. But most elements (C, N, F, s-elements) probably come from red giant stars • 12C from trippel , • 14N from C N O • s-elements from 13C+16O + n • Problem to get them up: He shell flashes! Promote mixing in episodes up to the deep convective envelope

  22. And then to get it up from the star! Mass loss observed -- for stars high up on AGB of 10-5 Msun/year or even more! How does it work? Pulsations, dust formation, radiative pressure in interaction Höfner and Freytag (2008)

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