Supernova Evolution and Explosive Nucleosynthesis

Transcription

1 Supernova Evolution and Explosive Nucleosynthesis Gabriel Martínez Pinedo 5 th European Summer School On Experimental Nuclear Astrophysics Santa Tecla, September 25, 2009

2 Outline 1 Introduction 2 Electron capture Presupernova Evolution Collapse phase 3 Explosive nucleosynthesis of heavy elements Nucleosynthesis processes Supernova neutrino heated ejecta Proton-rich ejecta: The νp-process

3 Outline 1 Introduction 2 Electron capture Presupernova Evolution Collapse phase 3 Explosive nucleosynthesis of heavy elements Nucleosynthesis processes Supernova neutrino heated ejecta Proton-rich ejecta: The νp-process

4 Outline 1 Introduction 2 Electron capture Presupernova Evolution Collapse phase 3 Explosive nucleosynthesis of heavy elements Nucleosynthesis processes Supernova neutrino heated ejecta Proton-rich ejecta: The νp-process

5 SN1987A Type II supernova in LMC ( 55 kpc) neutrinos E ν erg 50.0 Kamiokande II 40.0 IMB Energy (MeV) Time (seconds) light curve E grav erg E rad erg E kin erg = 1 Bethe

6 Stellar Evolution

7 Stellar life Nuclear burning stages (e.g., 20 solar mass star) Fuel Main Product Secondary Product T (10 9 K) Time (yr) Main Reaction H He 14 N CNO 4 H 4 He He O, C 18 O, 22 Ne s-process He 4 12 C 12 C) 16 O C Ne, Mg Na C + 12 C Ne O, Mg Al, P Ne) 16 O 20 Ne) 24 Mg O Si Si, S Fe Cl, Ar, K, Ca Ti, V, Cr, Mn, Co, Ni O + 16 O 28 Si)

8 Presupernova Star Star has an onion like structure. Iron is the final product of the different burning processes. As the mass of the iron core grows it becomes unstable and collapses when it reaches around 1.4 M. E b /A 40 Ca 56 Fe 86 Kr 107 Ag U 9 24 Mg 8 16 O I 174 Yb 208 Pb 12 C 7 He 6 6 Li He 2 1 H 0 H A

9 Schematical Evolution H 13 ~10 cm He O/Si Progenitor (~ 15 M ) 7 (Lifetime: 1 ~ 2 10 y) Fe ν ν 6 10 cm ν Dense Core ν ν M ν ~0.1 1 Sec. ν Late Protoneutron Star (R ~ 20 km) ν M Supernova 7 10 cm ν ν Sphere Hot Extended Mantle Dense Core Early "Protoneutron" Star M ν Shock ν M M ν ν e e + p n + νe and Photodisintegration of Fe Nuclei ν e cm "White Dwarf" (Fe Core) ~1 Sec. Collapse of Core (~1.5 M ) km/s (R ~ km)

10 Abundance Low energy astrophysical processes conserve number of nucleons: n = n i A i i n number nucleons per cm 3, n ρ m u = ρn A (CGS) n i number nuclei species i Abundance: Y i = n i n n i = ρn A Y i (changes in density are factored out) Mass fraction: X i = n im i ρ = n ia i m u ρ = Y i A i From conservation of number of nucleons: i Y i A i = i X i = 1

11 Example: Solar Abundances The Solar Abundances constitute the reference point for any nucleosynthesis study. Abundance relative to Silicon = BBN H He D C O Ne SiS Li B Be He burning C burning O burning H burning Ca Si burning Fe Ni Nuclear Statistical Equilibrium Ge Sr r s Neutron capture processes: s process r process Xe Ba r s Mass Number Pt Pb

12 Electron Abundance From charge neutrality: n e = n i Z i = n Y i Z i i i Introducing: Y e = n e n Y e = Y i Z i In general one can not define a lepton abundance. Lepton number is not locally conserved (neutrinos leave the system).

13 Nuclear Statistical Equilibrium Processes mediated by the strong and electromagnetic interactions are in equilibrium. Equilibrium between creation and destruction of nuclei. A(Z, N) Z p + N n + γ s µ(z, A) = Zµ p + (A Z)µ N Composition depends only on (T, ρ, Y e ). Y e is a parameter that is determined by weak processes that need to be explicitly considered.

14 Nuclear abundances in NSE Nuclei follow Boltzmann statistics: ( µ(z, A) = m(z, A)c 2 n(z, A) 2π 2 ) 3/2 + kt ln G(Z, A) m(z, A)kT with G(Z, A) the partition function: G(Z, A) = (2J i + 1)e Ei/kT π 6akT exp(akt), Results in Saha equation: i with a A/9 MeV 1 Y(Z, A) = G Z,A(T)A 3/2 2 A ( ρ m u ) A 1 ( ) 2π Yp Z Yn A Z 2 3(A 1)/2 e B(Z,A)/kT m u kt Composition depends on two parameters: Y p, Y n. They are determined from mass and charge conservation: Y i A i = 1; Y i Z i = Y e i i

15 Abundances in Nuclear Statistical Equilibrium 45 T= 9.01 GK, = 6.80e+09 g/cm 3, Y e = Z (Proton Number) Log (Mass Fraction) N (Neutron Number) 45 T= GK, = 3.39e+11 g/cm 3, Y e = Z (Proton Number) Log (Mass Fraction) N (Neutron Number)

16 Early iron core The core is made of heavy nuclei (iron-mass range A = 45 65) and electrons. Composition given by Nuclear Statistical Equilibrium. There are Y e electrons per nucleon. The mass of the core M c is determined by the nucleons. There is no nuclear energy generation which adds to the pressure. Thus, the pressure is mainly due to the degenerate electrons (relativistic), with a small correction from the electrostatic interaction between electrons and nuclei. As long as M c < M ch = 1.44(2Y e ) 2 M, the core can be stabilized by the degeneracy pressure of the electrons.

17 Onset of collapse There are two processes that make the situation unstable: 1 Silicon burning is continuing in a shell around the iron core. This adds mass to the iron core increasing M c. 2 Electrons can be captured by protons (free or in nuclei): e + A(Z, N) A(Z 1, N + 1) + ν e. This reduces the pressure and keep the core cold, as the neutrinos leave. The net effect is a reduction of Y e and consequently of the Chandrasekhar mass (M ch ) 3 Neutrinos produced by weak interactions can leave the star unhindered. They carry away energy and help to cool the collapsing core.

18 Initial conditions The dominant contribution to the pressure comes from the electrons. They are degenerate and relativistic: P n e µ e ρy e µ e µ e is the chemical potential of the electrons: µ e 1.11(ρ 7 Y e ) 1/3 MeV For ρ 7 = 1 (ρ = 10 7 g cm 3 ) the chemical potential is 1 MeV, reaching the nuclear energy scale. At this point is energetically favorable to capture electrons by nuclei.

19 How to determine the evolution Composition determined by NSE, function of temperature, density and Y e. Weak interactions are not in equilibrium. Change of Y e has to be computed explicitly (Y i = n i /n): Y e = Y i Z i Ẏ e = λ i ecy i + λ i β Y i i i i

20 Presupernova evolution R [km] R ~ 3000 Fe Initial Phase of Collapse (t ~ 0) Fe, Ni ν e Si ν e ~ M M(r) [M ] Ch Si burning shell ν e T = MeV, ρ = g cm 3. Composition of iron group nuclei. Important processes: electron capture: e +(N, Z) (N +1, Z 1)+ν e β decay: (N, Z) (N 1, Z+1)+e + ν e Dominated by allowed transitions (Fermi and Gamow-Teller) Evolution decreases number of electrons (Y e ) and Chandrasekar mass (M ch 1.4(2Y e ) 2 M )

21 Example of beta-decay: ft-value Q EC = (1) (1) fs 1.24% ~0.13% ~0.3% ~6.3 ~ (1) ps ps ps ps ps 26.8% 2.64% 0.21% 25.3% 5.9% 3.59% ps <0.25% > ns <0.6% >7.5 ps 1.80 ps ~ ~ < < D E M1(+E2) M1+E E E E E E E2+M E m Ga stable 0+ f t 1/2 = 64 30Zn K B(F) + B(GT), % 6.5 K = ± 2.4 s f (Q) phase space function. B(F) Fermi matrix element. B(GT) Gamow-Teller matrix element.

22 Gamow-Teller Transitions B(GT) = g 2 A 2J i + 1 J f ; T f T z f A σ k t± J k i ; T i T zi 2 k=1 g A = ± Sum rule (sum over all the final states): S = f O i 2 = i O O i f For Gamow-Teller transitions we have the Ikeda sum rule: S (GT) = S (GT) S + (GT) = 3(N Z) For a neutron, S (GT) = 3, S + (GT) = 0.

23 Systematics Electron capture 15 Beta decay 20 E (MeV) (Z 1,A) GT + GT F E (MeV) (Z,A) (Z 1,A) In neutron rich nuclei GT + strength represent a small part of Ikeda sum rule (S = 3(N Z) + S + so S S + ). In addition, its value is rather sensitive to many body correlations. For neutron rich nuclei GT constitutes most of the Ikeda sum rule. Most of the strength is located in a resonance with a width of several MeV and at energy: E GT E IAS = (N Z)/A MeV. Fermi transitions only appear in the β direction. All the strength (N Z) is located at the IAS state at an energy with respect of the parent state: Q IAS = E C (m n m p ) 0

24 Laboratory vs. stellar electron capture Laboratory Low lying Strength σ τ (Z 1,A) (Z,A) (Z 1,A) Supernova Gamow Teller Resonance (Z,A) electron distribution Capture of K-shell electrons to tail of GT strength distribution. Parent nucleus in the ground state Capture of electrons from the high energy tail of the FD distribution. Capture to states with large GT matrix elements (GT resonance). Thermal ensemble of initial states.

25 Beta-decay Electron capture 15 Beta decay 20 E (MeV) Finite T (Z,A) GT 5 F GT + 0 (Z+1,A) E (MeV) GT + and GT sum rules related by Ikeda sum rule: S S + = 3(N Z)

26 GT in charge exchange reactions GT strength could be measured in Charge-Exchange reactions: GT proved in (p, n), ( 3 He, t). GT + proved in (n, p), (t, 3 He), (d, 2 He). Mathematical relationship (E p 100 MeV/nucleon): dσ dωde (0 ) f (E x )B(GT) B(GT) = g 2 A 2J i + 1 f k σ k t k ± i 2

27 Independent Particle Model GT + strength in 58 Ni measured in (n, p) Ni (El-Kateb et al, 1994) Independent particle model. p 1/2 f 5/2 GT Strength FFN στ + f 5/2 p 1/2 p 3/2 f 7/2 ν π p 3/2 f 7/2 ν π f 5/2 p 1/2 p 3/ E (MeV) The IPM allows for a single transition ( f 7/2 f 5/2 ). It does not correctly reproduce the fragmentation of GT strength (correlations). To account for correlations, it is necessary to explicitly consider the residual interaction between nucleons. 58 Ni f 7/2 ν 58 Co π

28 Comparison with measured GT+ strength Nucleus Uncorrelated Correlated Expt. (IPM) Unquenched Q = V ± Fe ± Mn ± Fe ± Ni ± Co ± Ni ± Ne ± 0.2

29 Gamow-Teller strength and rates GT Strength Fe 56 Fe 58 Ni FFN Exp SM E (MeV) V 55 Mn 59 Co E (MeV) λ ec (s 1 ) Fe Ni LMP FFN ρ 7 = ρ 7 =10.7 ρ 7 = T Co 59 Co 54 Mn ρ 7 = Co ρ 7 =4.32 ρ 7 = T 9 G2 F V2 ud σ i, f (E e ) = F(Z, E 2π 4 c 3 e )B i, f (GT)Eν 2 v e λ i, ec f = 1 p 2 π 2 3 e f (E e, T, µ e )σ i, f (E e )de e Q i f λ ec = i, f (2J i + 1)e E i/kt λ i, f ec i(2j i + 1)e E i/kt

30 KVI results using (d, 2 He) C. Bäumer et al. PRC 68, (2003) B(GT + ) V(d, 2 He) 51 Ti λ (s 1 ) LMP (d, 2 He) T (10 9 K) Old (n, p) data V(n,p) Alford et al. (1993) 0.4 B(GT + ) 0.3 large shell model calculation E x [MeV] B(GT+) GT strength in 48 Sc, 50 V, 58 Ni, 64 Ni also measured.

31 NSCL results using (t, 3 He) A. L. Cole, et al., PRC 74, (2006) λec (s 1 ) (d, 2 He) KB3G LMP FFN GXPF λec/λec,exp T (10 9 K) Extension to unstable nuclei requires measurements in inverse kinematics

32 Consequences weak rates A. Heger et al., PRL 86, 1678 (2001) T (10 9 K) Ye M 25 M T ρ WW LMP T ρ ρ (g cm 3 ) 0.44 (s 1 ) dy e dt LMP EC LMP β Time till collapse (s) Time till collapse (s) 10 0

33 Ye evolution vs beta equilibrium How does the Y e evolution compares with the assumption of beta equilibrium (µ e + µ p = µ n ). Beta equilibrium defines the lowest Y e that can be achieved during evolution. Figure shows evolution 15 M star. Ye equilibrium LMP Woosley & Weaver time [s]

34 Important variables end presupernova A. Heger et al., ApJ 560, 307 (2001) Ye WW LMP MFe (M ) WW LMP central entropy / baryon (k B ) WW LMP Ye MFe (M ) S (k B ) Star Mass (M ) Star Mass (M ) Star Mass (M )

35 Collapse phase R [km] R Fe ~ heavy nuclei Neutrino Trapping (t ~ 0.1s, c~10¹² g/cm³) δ Fe, Ni M hc 1.0 ν e Si ν e ν e M(r) [M ] Si burning shell Important processes: Neutrino transport (Boltzmann equation): ν + A ν + A (trapping) ν + e ν + e (thermalization) cross sections E 2 ν electron capture on protons: e + p n + ν e electron capture on nuclei: e +A(Z, N) A(Z 1, N+1)+ν e

36 (Un)blocking electron capture at N=40 Independent particle treatment g 9/2 N=40 Blocked GT f 5/2 p 1/2 Unblocked Correlations Finite T g 9/2 f 5/2 p GT 1/2 p 3/2 f 7/2 neutrons protons Core p 3/2 f 7/2 neutrons protons Core

37 Experimental information for 76 Se and 76 Ge 76 Se (Z = 34, N = 42) has been the subject of several recent studies due to the interest in the double-beta decay of 76 Ge (Z = 32, N = 44). 76 Se(d, 2 He), Grewe et al, PRC 78, Measured occupation numbers (2008) Schiffer et al, PRL 100, (2008); Kay et 1 H 12 B (g.s.) al arxiv: Se p f 5/2 g 9/2 neutrons ( 1.59) ( 2.17) (+3.80) protons ( 1.92 ) (+1.16) (+0.84) 76 Ge p f 5/2 g 9/2 neutrons ( 1.13) ( 1.44) (+2.48) protons ( 2.23) (+2.04) (+0.23) ΣB(GT; Se76->As76) MeV (1 +,2-) 0.5 ± 0.3 MeV (1 +,2-) 1.03 MeV (1 + ) 1.22 MeV (2 -) 1.63 MeV (1 +) 1.86 MeV (1 +) 2.22 MeV (1 +) 2.4 MeV (2 +) E x [MeV] K. Sieja (Shell Model calculation) SM EXP E (MeV)

38 Electron capture: nuclei vs protons Electron capture rates 40 Energetics λec (s 1 ) H 68 Ni 69 Ni 76 Ga 79 Ge 89 Br (MeV) µ e ρ (g cm 3 ) Abundances Y n 10 0 Q = µ n µ p Q p ρ (g cm 3 ) Abundance Y h Y α Y p R h = i Y i λ i = Y h λ h R p = Y p λ p, Y i = n i /n ρ (g cm 3 )

39 Effects Realistic calculation With Marek, Rampp, Janka & Buras (Approx. General Relativistic model) 15 M presupernova model from A. Heger & S. Woosley Ye,c, Ylep,c Bruenn LMSH Y e Y lep ρ c (g cm 3 ) Eν emission (MeV s 1 ) Bruenn LMSH Eν (MeV) ρc (g cm 3 ) ρ c (g cm 3 ) Ye Ylep ρ [g/cm 3 ] Electron capture on nuclei dominates over capture on protons All models converge to a norm stellar core at the moment of shock formation. s11-slms s15-slms s20-slms s25-slms Ye

40 4 Implications for the collapse Improved treatment of electron capture produces: Reduced core size at bounce. Reduction in the average energy of the emitted neutrinos. Y e (Electron Fraction) Velocity (10 km/s) Bruenn prescription LMSH rates Enclosed Mass Luminosity (10 46 W) Neutrino Energy (MeV) Bruenn prescription: e neutrino e antineutrino LMSH rates: e neutrino e antineutrino Time relative to bounce (s) W. R. Hix et al, PRL 91, (2003)

41 Nucleosynthesis beyond iron Three processes contribute to the nucleosynthesis beyond iron: s-process, r-process and p-process (γ-process) Z p-drip n-drip N Abundances [Si=10 6 ] r s r s 115 Sn 180W p 138 La 152 Gd 164 Er 180 m Ta A s-process: relatively low neutron densities, τ n > τ β r-process: large neutron densities, τ n < τ β. p-process: photodissociation of s-process material.

42 BBFH suggested that neutron deficient nuclei are produced by proton captures in an evironment with proton densities of 10 2 g cm 3 and T 2 3 GK. Can we find an astrophysical scenario with proton-rich ejecta? Nucleosynthesis beyond iron p-process fails to explain the solar abundances of 92,94 Mo and 96,98 Ru [Arnould & Goriely, Phys. Rep. 384 (2003) 1]

43 r-process and metal-poor stars Cowan & Sneden, Nature 440, 1151 (2006) a log ε(x) b log ε(x) Ga Ge 1 30 Sr Y Zr Mo Nb Ru Pd CdSn Rh Ag Solar System s-process only Solar System r-process only CS detections CS upper limits Os Pt Pb Ba Nd Sm DyEr Ir Ce Gd Yb La Pr Eu Ho Tb Normalized at Eu Tm Lu Atomic number Hf Au Observed minus Solar System r-process only Th U Nucleosynthesis Signatures of few (single?) events. Abundances Z > 56 consistent with solar r-process abundance. Z < 56 are underproduced with respect solar abundances. At least two different astrophysical sites are necessary to explain solar r-process abundances. What is the origin of the robust r-process?

44 r-process and metal-poor stars Cowan & Sneden, Nature 440, 1151 (2006) Relative log ε CS HD BD Atomic number CS HD SS r-process abundances Several stars show robust r-process for Z > 56. Nucleosynthesis Signatures of few (single?) events. Abundances Z > 56 consistent with solar r-process abundance. Z < 56 are underproduced with respect solar abundances. At least two different astrophysical sites are necessary to explain solar r-process abundances. What is the origin of the robust r-process?

45 r-process generalities The r-process is a primary process that requires very high neutron densities. Astrophysical scenarios may be classified in two categories: High entropy environments (fast expanding, low Y e ejecta). High temperature r-process (classical (n, γ) (γ, n) equilibrium) Low temperature r-process (competition between neutron captures and beta decay, Blake & Schramm, ApJ 209, 846 (1976) Low entropy environments (ejection cold neutron-rich matter, low temperature r-process). Both scenarios differ in the way neutrons and seed nuclei are created but not so much in the way the r-process operates.

46 Weak Interaction in Explosive Nucleosynthesis ν ν ν ν ν ν ν ν ν ν ν ν ν ν ν Main processes: ν e + n p + e ν e + p n + e + Neutrino interactions determine the proton to neutron ratio. o νp-process (proton-rich ejecta) r-process (neutron-rich ejecta)

47 Nucleosynthesis in neutrino heated matter Proton rich (νp-process) νe + p n + e + 64 Ge + n 64 Ga + p 64 Ga + p 65 Ge;... seeds (N = Z 28 32) Heating region... 4 He(αα, γ) 12 C 2p + 2n 4 He Alpha formation Weak interaction freeze out νe + n p + e νe + p n + e + Proto neutron Star Neutron rich (r-process) neutrons + seeds heavy nuclei (A 100??) seeds (A )... 4 He(αα, γ) 12 C 4 He(αn, γ) 9 Be Heating region { 0.25 MeV T 3 GK { 0.75 MeV T 9 GK { 0.9 MeV T 10 GK νe, νe, νµ, νµ, ντ, ντ Matter is ejected under the influence of strong neutrino fluxes. Nucleosynthesis sensitive to entropy and expansion time scale that depend in neutrino properties (Luminosity and mean neutrino energy) and proto neutron star properties (Mass and radius) [Qian & Woosley, 1996]. Composition (Y e ) depends in spectral differences between neutrinos and antineutrinos. After alpha rich freeze-out nucleosynthesis depends of the proton to seed ration (proton rich) or neutron to seed ration (neutron rich).

48 Why proton or neutron rich? Y e evolution is mainly determined by the competition between ν e + n p + e and ν e + p n + e + : dy e = λ νe ny n λ νe py p = λ νe n (λ νe n + λ νe p)y e, dt using Y e = Y p and Y n = 1 Y n. If matter is exposed long time enough to neutrino fluxes it will reach an equilibrium given by: Y e = λ νe n. λ νe n + λ νe p We need to estimate the neutrino absorption rates?

49 Why proton or neutron rich? If we assume that the produced electron (positron) is extremelly relativistic (E e = p e c) the cross section for (anti)neutrino absorption is: σ(e ν ) = σ 0 (E ν ± ) 2 plus sign for neutrinos and minus sign for antineutrinos, = m n c 2 m p c 2 = MeV, and σ 0 = (1 + 3g2 A )G2 F V2 ud π 4 c 4 = 9.33(4) cm 2 MeV 2 For the absorption rate we need to integrate over the neutrino spectrum and multiply by the neutrino flux (L ν /(4πr 2 E ν ) obtaining: λ νn = L ) ν 4πr 2 σ 0 (ɛ ν E ν ) λ νp = L ν 4πr 2 σ 0 (ɛ ν E ν with ɛ ν = E 2 ν / E ν.

50 Why proton or neutron rich? Now we can assume that L ν L ν and we find that in order to achieve Y e > 0.5 we need: Y e = λ νe n λ νe n + λ νe p > 1/2 ɛ ν ɛ ν < 4 That means that if the difference between antineutrino and neutrino energies is smaller than 5.2 MeV we can have proton rich ejecta even if antineutrinos have more energy than neutrinos.

51 Evolution of neutrino fluxes and energies: Y e Luminosity [10 51 erg/s] Luminosity [10 51 erg/s] rms Energy [MeV] Fischer et al, arxiv: , have performed the first Boltzmann neutrino transport study of the proto-neutron star evolution (a) 10 solar mass Time After Bounce [s] (b) ν e anti ν e ν µ/τ anti ν µ/τ Time After Bounce [s] (c) ν e anti ν e ν µ/τ Luminosity [10 51 erg/s] Luminosity [10 51 erg/s] rms Energy [MeV] (a) 18 solar mass Time After Bounce [s] (b) ν e anti ν e ν µ/τ anti ν µ/τ Time After Bounce [s] (c) ν e anti ν e ν µ/τ Electron Fraction, Y e Ejecta is always proton-rich Boltzmann ~(L ν,< ε ν >) ~(λ ν, λ anti ν ) Time after bounce [s] Nucleosynthesis is sensitive to the interaction of wind with previously ejected matter (reverse shock) [A. Arcones et al, A&A 467, 1227 (2007)] Time After Bounce [s] Time After Bounce [s]

52 α-process: Seed build up Typically three-body reactions are necessary to build up seed nuclei from alpha particles ( ρ 2 ): 4 He(αα, γ) 12 C is the dominating process in proton-rich ejecta. 4 He(αn, γ) 9 Be(α, n) 12 C is the dominating process in neutron-rich ejecta. In neutron rich ejecta the sequence of reactions 4 He(t, γ) 7 Li(n, γ) 8 Li(α, n) 11 B can enhance the production of seed nuclei (Sasaqui et al, 2006; Terasawa et al, 2001). In general a successful r-process requires a large neutron to seed ratio (> 100). This demands low Y e (large amounts of neutrons), fast expansions (milliseconds) and/or large entropies (> 150 k B small build up of seed nuclei).

53 Adiabatic expansion studies Y e τ dyn = s s s S Combinations Y e, s, and τ dyn necessary for producing the A = 195 peak. Main problem α effect, r-process in a strong neutrino flux. We start with Y p protons and Y n neutrons (Y n Y p ) Alpha particles form (Y α = Y p, Y n = Y n Y p ) New protons created by neutrino interactions Final results is a reduction of neutron abundance (n/seed) that makes imposible r-process. Possible solutions: very fast expansion, neutrino oscillations,...

54 Iron group nuclei and proton rich ejecta Proton-rich ejecta are the mayor contributors to 45 Sc and 64 Zn (C. Fröhlich, et al. 2006, J. Pruet, et al. 2005) [X/Fe] Cayrel et al. Gratton & Sneden Without proton-rich With proton-rich Ca Sc Ti V Cr Mn Fe Co Ni Cu Zn Model and observations in metal-poor stars.

55 Impact neutrino interactions in proton-rich ejecta Neutrino interactions are responsible for the production of nuclei with A > Without ν With ν Yi/Yi, Mass Number A

56 The νp process Proton rich matter is ejected under the influence of neutrino interactions. Nuclei form (mainly N = Z) at distances where a substantial antineutrino flux is present. Antineutrino charge-current capture time and expansion time scale are similar ( 1 s) Neutrinos speed-up matter flow ν e + p e + + n n + 64 Ge 64 Ga + p 64 Ga + p 65 Ge As (p, γ) S 74 kev p 64 Ge β s 64 Ga These reactions constitute the νp-process C. Fröhlich, et al., PRL 96, (2006)

57 The νp process Proton rich matter is ejected under the influence of neutrino interactions. Nuclei form (mainly N = Z) at distances where a substantial antineutrino flux is present. Antineutrino charge-current capture time and expansion time scale are similar ( 1 s) Neutrinos speed-up matter flow ν e + p e + + n n + 64 Ge 64 Ga + p 64 Ga + p 65 Ge... These reactions constitute the νp-process C. Fröhlich, et al., PRL 96, (2006) 65 As 64 Ge ( n,p) ~ 1 s 66 As 65 Ge 64 Ga (p, γ) (p, γ)

58 The νp process Supernova model courtesy from H.-Th. Janka Proton rich matter is ejected under the influence of neutrino interactions. Nuclei form (mainly N = Z) at distances where a substantial antineutrino flux is present. Antineutrino charge-current capture time and expansion time scale are similar ( 1 s) Neutrinos speed-up matter flow ν e + p e + + n n + 64 Ge 64 Ga + p 64 Ga + p 65 Ge... These reactions constitute the νp-process C. Fröhlich, et al., PRL 96, (2006) Mi/(M ej Xi, ) Cu Se Kr Zn GaGe Br Rb As Sr Mass number A Y Mo Zr Nb Sensitivity to masses and astrophysical conditions. Ru Pd Rh Cd

59 Evolution some abundances log10(yi) Ni n p 4 He 1 10 Time (s) During the early phases proton-rich ejecta are in NSE. Differently what happens in neutron-rich ejecta, proton-rich NSE has a composition that is basically given by N = Z nuclei, alphas and free protons (Seitenzahl et al, arxiv: [astro-ph]) Y e = Y p + (X α + X h )/2, Y p + X α + X h = 1, Y p = 2Y e 1

60 Estimation of neutron abundance Neutron abundance is determined by an equilibrium between neutron production Y p λ νe p and destruction i Y n Y i ρn A σv i (n,p). Taking typical values for the different quantities (T = 2 GK, Y e = 0.57): Y p = 0.14; λ νe p = 0.1 s 1 N A σv (n,p) = cm 3 mol 1 s Y h = 10 4 ; ρ = 10 4 g/cm 3 we get Y n This results in an averate destruction rate of a nucleus by (n, p) reactions of Y n ρn A σv (n,p) = 100 s 1.

61 Masses measured at SHIPTRAP and JYLFTRAP ν C. Weber et al, Phys. Rev. C 78, (2008)

62 Influence new masses Mi/(M ej Xi, ) Ratio Zn Se Ge Ga Kr Sr As Br Se Kr Sr Mo Rb Zr Y Nb Ru Mo Ru Pd Pd Cd AW Expt Mass number New set of reaction rates using new experimental masses (T. Rauscher). New experimentar values close to Audi-Wapstra systematics. Little changes in abundances. No spectroscopic information is known for these nuclei. Sensitivity studies of the important reaction rates are necessary. However, νp-process seems to occur under (p, γ) (γ, p) equilibrium. Having the nuclear physics under control allows to explore the astrophysical uncertainties. The νp-process can be used constrain the evolution of the ejected matter.

63 Sensitivity to temperature evolution The νp-process nucleosynthesis is rather sensitive to the temperature evolution of the ejected matter. In particular, to the interaction with the reverse shock. Temperature (GK) Time (s) (A. Arcones & GMP, in preparation) Mi/(M ej Xi, ) Cu Kr Se Sr Ga Ge Br Rb Zn As Zr Nb Y Trs = 2 GK Trs = 3 GK Trs = 1 GK no rs Mass number A Mo Ru Pd Rh Cd

64 Production of light p nuclei S. Wanajo, astro-ph/ Mo is produced in slightly neutron rich ejecta, Y e (Fuller & Meyer 1995, Hoffman et al. 1996). The rest is produced in proton-rich ejecta.