Stellar Abundances
As probes of stellar nucleosynthesis and chemical evolution
I. Basic principles of stellar nucleosynthesis II. Basic ingredients III. Deciphering abundances
Francesca Primas ESO KES Lecture – Mar 6, 2017 GCE: some basics
The chemical history of the Galaxy is dominated by the nucleosynthesis occurring in many generations of stars
• Ini ally: X=Xi (~0.75) Y=Yi (~0.25) Z=Zi (0.0) Today: X=Xj (~0.73) Y=Yj (~0.25) Z=Zj (0.02)
• Elements are made in stars, ejected into the ISM by stellar winds and supernova explosions.
• Elements are mixed into the ISM out of which new stars form. Each new stellar genera on is born from gas with higher metallicity:
X Y Z yield, IMF, SFR
• The cycle repeats itself over billions of years
Chemical abundances in the solar system/ vicinity provide the richest informa on. Different ways to sort and group the elements …
• By nucleosynthe c origin and nuclear proper es • primordial, H-, He-, C-, Ne-, Si-burning, n-captures • stable, long-lived or short-lived radioac ve
Big Bang CR spalla on H-burning (hydro) H-burning (expl) He-burning C-burning O-burning Ne-burning Expl. Nucleosyn. NSE s-process r-process Origin of atoms in the solar system
• Two sources: nucleosynthesis in the BIG BANG and in STARS
• The Big Bang made only H and He (and some Li)
• All other nuclei are created/produced in stars, by 3 kinds of processes:
Nuclear burning (fusion) PP cycles, CNO bi-cycle, He-, C-, O-, Si-burnings … make atoms up to 40Ca (in MS + RG stars)
Photodisintegra on rapid nuclear rearrangements, at γ-ray energies 56 ESOTHERMIC creates everything up to Fe
Neutron irradia on nuclei heavier than 56Fe are generated by n-captures, following two paths (n-flux):
à The s-process, in which neutron addi on is slow compared to β-decay à The r-process, in which neutron addi on is rapid compared to β-decay (only in SN) 4
ENDOTHERMIC Proton irradia on similar to s-process but with protons, makes low-abundance nuclei The 1957 milestone B2FH – 1957
• Hydrogen burning: responsible for most energy produc on in stars – all cycles synthesizing He from H + isotopes of C,N,O,F,Ne and Na (not produced in He+α)
• Helium burning: responsible for synthesis of C from He + produc on of O16, Ne20 and maybe Mg24 with extra α-s
• The α process: adding α par cles to Ne20 to form Mg24, Si28, S32, A36, Ca40 (and probably Ca44 and Ti48)
• The e process: the equilibrium process (very high T and ρ) makes the iron-group (V,Cr,Mn,Fe,Co,Ni)
• The s process: n-capture with emission of (n,γ) on a long mescale (100yrs-105yrs/n- capture); 23 • The r process: n-capture on short mescales (0.01-10s); large frac on 70 • The p process: p-capture with emission of (p,γ) or (γ,n); responsible for p-rich isotopes, with very low abundances • The x process: synthesis of D,Li,Be,B (unstable at Tint, thus requiring regions of low T and ρ) … and the νp process B2FH: Atomic abundances curve as a function of atomic weight Notes from observa on of atomic absorp on • H most abundant, then general lines in the solar spectrum decrease ll U (least abundant) • Big nega ve anomaly at Be, B, Li • Moderate posi ve anomaly around Fe • Sawtooth pa ern from odd-even effect Successful model of nuclear origins needs to explain all these features ! in the abundance pattern!! Hydrostatic H-burning Sun (T=15.6 MK), stellar core (T=8-55 MK), shell of AGB stars (T=45-100 MK) • At the starts, none of the 2-par cle reac ons between H and He is stable • 1H + 1H = 2He (unstable) = 1H + 1H • 1H + 4He = 5Li (unstable) = 1H + 4He • 4He + 4He = 8Be (unstable) = 4He + 4He • 1939: Hans Bethe shows how H-burning can begin with the exothermic forma on of D: • 1H + 1H = 2D + β+ + ν + 1.442 MeV • This reac on ini ates the PP-I chain: 2 (1H + 1H = 2D + β+ + ν) 2 ( 1H + 2D = 3He + γ) 3He + 3He = 4He + 2 1Η 1 4 Net: 4 H = He + 2 ν + γ + 26.7 MeV 8 2D/1H quickly approaches equilibrium value, but 1013 times << than the terrestrial value… PP-chain: main reactions, importance, lifetimes 90% of Sun’s energy produced by pp1 chain The rate at which reac ons occur depends on the par cle energy, flux and the reac on σ + CNO-cycle, He, C, Ne, O, Si burnings … Nuclear binding energy 56 Neutron capture Fe reactions can proceed As soon as we have neutrons, they are captured by nuclei Fusion exothermic α to their n- capture σ A Evidence: abundance correla on In principle, nuclear burning by with n-capture σ fusion can continue only up to 56Fe, the nucleus with the greatest binding energy per nucleon. 10 H-burning is by far the most effective means of converting mass into energy! Neutron capture processes • If neutron flux is slow compared to β-decay mes, nuclei follow the valley of stability and make s-process nuclei • If neutron flux is so fast that repeated captures occur before β-decay, nuclei on the neutron dripline (where the capture rate goes to zero) are made, which subsequently decay back to first stable nuclide on each isobar. These make the r-process (r = rapid) • These require energy and occur only at high T and ρ: Core+shell burning p- & s-process Supernovae p- & r-process r-process departs from stability valley 11 s-process After Si exhaustion in the core … No fuel le … core grows in mass – At MC > 1.4Msun : e degeneracy pressure ≠ gravity Core collapses At ρ=1014 g/cm3 : part of core rebounds à outward moving shock wave Composition of layers dominated by more stable nuclei (A multiple of 4) Explosive Si-burning (T=4-5 GK) High T and ρ reached in inner 28Si layer because of outgoing shock wave Ma er cools when shock moves outwards NSE à non-NSE If because of lack of α-par cles à “α-par cle-poor freeze-out” à ejected abundances ≈ NSE (mainly 56Ni) Nucleosynthesis now depends critically on ρ, If because of excess of α-par cles à “α-rich freeze-out” expansion time scale à ejected abundances somewhat ≠ NSE (s ll 56Ni but also 44Ti) and n, p, α abundances Explosive O-burning (T=3-4 GK) Nucleosynthesis similar to hydrosta c O-burning Main source of 28Si, 32S, 36Ar, 40Ca (“α-elements”) in Universe Explosive C- and Ne-burnings (T=2-3 GK) Nucleosynthesis similar to hydrosta c C- and Ne-burnings Predicted to be main source of 26Al Stellar nucleosynthesis: a recap Source: Alex Heger 8 M/Msun Fuel Products T/10 K 0.08 H He 0.2 1.0 He C, O 2 Not all stars undergo the AGB complete sequence à mass! 1.4 C O, Ne, Na 8 5 Ne O, Mg 15 Only massive stars go straight up to the end. 10 O Mg … S 20 20 Si Fe … 30 >8 SNe All! And now ? Source: Nomoto, Kobayashi & Tominaga 2013 Testing the different types Source: Nomoto, Kobayashi & Tominaga 2013 Stellar Abundances As probes of stellar nucleosynthesis and chemical evolution I. Basic principles of stellar nucleosynthesis II. Basic ingredients III. Deciphering abundances Francesca Primas ESO KES Lecture – Mar 6, 2017 Nomenclature Mass frac ons/ X=H; Y=He, Z=all other elements (=metals) Concentra ons X+Y+Z = 1 Xsun= 0.7381, Ysun= 0.2485, Zsun= 0.0134 (Asplund et al. 2009) The 12 scale log ε(X) = log (nX/ nH) + 12 (log ε(H) ≡12) E.g. log ε(O)sun≈8.7 dex (O is ≈2000× less abundant than Hsun) The [] scale [X/H] = log (nX/ nH)* – log (nX/ nH)sun E.g. [Fe/H] = –2 → star has 1/100 less iron than the Sun [Fe/H] = 0.5 → star has 3.16× more iron than the Sun How do you extract elemental abundances from these lines? • Star ID (spectral type or photometric classifica on) • Atmospheric proper es → line forma on • Effec ve temperature • Surface gravity • “Metallicity” Stellar photosphere: a de inition Transi on from interior to ISM Layer from which we receive photons Source: NASA Source: Bessell 2005 How to determine Teff Mul -colour photometry calibrated with stars having fundamentally determined Teff – Johnson’s UBV – Strömgren’s uvby – Geneva UBV B1 B2 V1 G Spectroscopy excita on equilibrium ra os of suitable strong lines → Teff or spectral type + calibra on Teff vs spectral type line profiles How to determine log g Calibrated photometry Parallaxes L, R, (M) → g à Gaia Spectroscopy Ra o of, e.g., Fe II to Fe I lines Line profiles: B-to-F-type stars: Balmer Late-type stars: strong metallic lines with wings (e.g. Na ID doublet, Ca II IR triplet) But there is always a temperature dependence – this must be secured beforehand or a pressure-temperature solution must be made. How to determine metallicity Lines iso-(B2-V1) and iso-m2 of the Geneva photometric system Calibrated photometry A er Teff (and maybe reddening) the global metallicity has largest influence on stellar flux Difficult for stars with [Fe/H] < –2 (e.g. δ(U –B) loose sensi vity) Limited precision: ≈0.3 dex 5000 6000 7000 8000 Teff Spectroscopy Equivalent widths From spectral lines to abundances Different methods and tools, but … Intensity/strength element in original stellar gas Find number Na of absorbing atoms per unit area Boltzmann and Saha equa ons are applied and combined with the pressure and Tgas to derive an abundance of the element (i.e. to calculate excita on and ioniza on): Saha à rela ve # atoms in 2 ioniza on states as a f(ne,T) Boltzmann à popula on of energy levels a,b Examples: metal lines Line strength depends on level popula on of the atoms. Strongest when temperature is low (lower ioniza on stages are populated) Dominate in F, G, K stars Source: Tolstoy Source: Subaru Press Release, Ito et al. 2009, ApJL, 698,37 Examples: dominant features Measuring abundances: equivalent width Equivalent width: Geometrically: Wλ =width of a rectangular, completely opaque line Curve of growth (EW varies with Na) Weak lines à linear part à W ∝ A Strong line Satura on Satura on à plateau à W ∝√ log A Weak-line Strong lines à damp. wings dominate à W ∝ √A The Solar Spectrum Source: Kurucz Crowding & Blends Spectrum Synthesis Stellar Abundances As probes of stellar nucleosynthesis and chemical evolution I. Basic principles of stellar nucleosynthesis II. Basic ingredients III. Deciphering abundances Francesca Primas ESO KES Lecture – Mar 6, 2017 GCE modeling Source: Kobayashi et al 2006 The [α/Fe] ratio The [α/Fe] ratio for real plateaus knees Halo Thick Thin Source: Venn et al. 2004, Shetrone 2013 Mg/Fe in the Local Group UMaII, ComberI, Draco, Umi Sculptor Fornax Carina Sextans Source: Jablonka 2013 From a 4m to an 8m telescope … Very inhomogeneous, single SNe? McWilliam et al (1995)! From a 4m to an 8m telescope … A rather well mixed early Galaxy!!! Very inhomogeneous, single SNe? Cayrel et al (2004) McWilliam et al (1995)! Bonifacio et al (2009)! Heavies Some very metal-poor stars are rela vely rich in neutron-capture elements, and their abundance pa ern indicates dominance of r-process synthesis. Source: Sneden et al. r-process rich EMP stars • 1st, 2nd, 3rd peak! Universal ? Nucleo-chronometry: • cosmo-chronometry (!) U/Th, U/r, Th/r Δt = k (log(X/Y)o - log(X/Y)obs) Light n-capture elements ! Eu, Os, Ir, Pt: clear positive correlations ! common nucleosynthesis origin! Ge: factors 4-5 lower! Either the r-process does not have anything to do with the creation of Ge, or these massive neutron blasts pushed all nuclei far beyond the Ge mass range. ! Lighter neutron-capture elements have different nucleosynthetic origin(s) than the elements with Z > 55 How robust are these results? Source: Gray Te End