Theory for nuclear processes in stars and nucleosynthesis

Transcription

1 Theory for nuclear processes in stars and nucleosynthesis Gabriel Martínez Pinedo Nuclear Astrophysics in Germany November 15-16, 2016 Nuclear Astrophysics Virtual Institute

2 Outline 1 Ab-initio description of nuclear reactions 2 Weak processes in stars 3 r process nucleosynthesis

3 Ab-initio description of nuclear reactions Extending ab-initio approaches for a proper treatment of continuum. Cluster phenomena in light nuclei (Hoyle state) Different approaches available. Fermion Molecular Dynamics No-Core Shell Model with Continuum Nuclear Lattice Dohet-Eraly, et al, PLB 757, 430 (2016)

4 Stellar electron capture and beta-decay We need to account for a broad range of conditions Stellar evolution and accretion phases. Sensitive to relatively few reactions: URCA pairs 23 Na- 23 Ne, 25 Mg- 25 Na,... for white dwarfs and heavier nuclei neutron star crust. Accurate modeling of individual transitions, important screening corrections. Explosive phases. Type Ia supernova, Oxygen deflagration in ONe cores. Competition between many electron capture and beta-decay processes. Core-collapse: core evolution sensitive to electron capture on exotic neutron rich nuclei.

5 Figure from Sam Jones Range of nuclei for Oxygen deflagration many of the relevant rates are based on simple analytical estimates (ANA).

6 Electron captures during collapse Most relevant electron capture nuclei during collapse e (b) Proton Number = N Z log 10 e Neutron Number Sullivan et al., ApJ 816, 44 (2015) Mainly nuclei around N = 50 and N = 82. Sensitive to shell structure far from stability. Theoretical challenge: description of correlations across shell closures. Many of the relevant nuclei are becoming experimentally accessible

7 Neutrino-matter interactions in supernova Y e Ref. run + weak magn. + wm + n decay S Y e S [k B /baryon] t t bounce [s] Many different processes contributing during different phases: explosion, neutron star deleptonization. Consistent treatment with underlying experimentally constrained EoS. Code implementations often introduce errors of order the claimed accuracy. What is their role in mergers?

8 Neutrino nucleus reactions Having accurate neutrino spectra is also important for the production of several nuclei ( 7 Li, 11 B, 19 F, including 26 Al). 26 Al (p,γ) 25 Mg (ν,ν'np) (ν,ν' n) (ν e,e - ) 28 Si 27 Al 26 Mg GT: ( 3 He, t), Zegers+ 2006; Forbidden: RPA cross section (10 42 cm 2 ) Mg + ν e 26 Al + e GT and Fermi transitions (data) Including forbidden transitions (RPA) Transitions to bound states, experimental Transitions to bound states Neutrino energy (MeV) low: E νe = 8.8 MeV, E νe = E νµ,τ = 12.6 MeV high: E νe = E νe = 12.6 MeV, E νµ,τ = 18.9 MeV M ) 26 Al yield ( KEPLER high energies low energies without ν Main sequence mass (M ) A. Sieverding E 2 ν f ν ( E ν =12 MeV)

9 Making Gold in Nature: r-process nucleosynthesis Solar r abundances N= Known mass Known half life r process waiting point (ETFSI Q) N=126 r process path N=184 The r-process requires the knowledge of the properties of extremely neutron-rich nuclei: Nuclear masses. Beta-decay half-lives. Neutron capture rates. Fission rates and yields.

10 r process in dynamical ejecta from mergers Ejection of very neutron-rich material in mergers results in abundance distributions insensitive to variations of astrophysical conditions. Sensitivity to nuclear physics input remains. abundances at 1 Gyr FRDM WS3 abundances at 1 Gyr HFB21 DZ mass number, A mass number, A

11 Theoretical masses far from stability Theoretical models do rather well far from stability! Sn (Z=50) Isotopes Mumpower Spread due to small differences in symmetry energy ( 0.5 MeV, smaller experimental range ).

12 Comparison S 2n Very similar predictions for Q-values (relevant quantity). S2n (MeV) FRDM DZ31 WS3 HFB 27 AME 2012 Sn (Z=50) Isotopes Neutron Number Variations in localized regions responsible for different abundances predictions.

13 Outlook: addressing systematic diff. between mass models Most of the differences between mass models originate due to: Treatment of transitional nuclei (shape coexistence). Requires beyond mean field techniques (Rodríguez+ 2015) Proper description odd and odd-odd nuclei. (3) (N) = (S n (N + 1) S n (N))/2 A. Arzhanov, Gogny functional 1.5 Sn (Z = 50) Isotopes Sn (3) (MeV) FRDM DZ31 WS3 HFB 27 AME Neutron Number (3) n [MeV] Pert. QP Pert. T Odd Equal Filling Full Blocking AME Number of Neutrons

14 ) - s 3 Constraining neutron capture rates Liddick, et al., Phys. Rev. Lett. 116, (2016) -1 (E x ) (MeV ) (a) -3 f(e ) (MeV ) (b) 1 N A < > (cm mol Ni(n, ) 70Ni rate (c) BRUSLIB JINA REACLIB E x (MeV) E (MeV) T(10 K) Experimental constrains in gamma-strength far from stability. Understanding low energy upbend and behaviour with neutron excess. Extending reaction model beyond statistical treatment.

15 Global beta-decay calculations Beta-decay rates determine the speed of matter flow from light to heavy nuclei. r-process path determined by neutron separation energies nuclei with largest impact are those with larger instantaneous half-lives. Despite tremendous progress at RIB facilities (RIBF at RIKEN) most of the half-lives are based on theoretical calculations. Two microscopic calculations (GT+FF) have become available: Covariant density functional theory + QRPA (Marketin+ 2016) Skyrme finite-amplitude method (Mustonen & Engel 2016) Marketin & GMP 2016 proton number log 10 T 1/2 D3C /T 1/2 F RDM neutron number Mustonen & Engel 2016

16 Fission barriers The impact of different fission barriers and yields has not been sufficiently explored. Goriely & GMP 2015 S. Giuliali Giuliani et al 2016 Goriely et al 2007 Möller et al 2009

17 Kilonova light curve Understanding the nuclear physics signatures in kilonova light curves L bol (ergs s 1 ) FRDM (Beta-decay dominates) DZ31 (Alpha-decay dominates) Days Ratio of luminosities at peak value and at late times can be used to constrain the produced amount of nuclei between Pb and U. Barnes, Kasen, Wu, GMP, ApJ 829, 110 (2016).