Nuclear physics input for the r-process

Transcription

1 Nuclear physics input for the r-process Gabriel Martínez Pinedo INT Workshop The r-process: status and challenges July 28 - August 1, 2014 Nuclear Astrophysics Virtual Institute

2 Outline 1 Introduction 2 Nucleosynthesis in supernova neutrino-driven winds 3 Nucleosynthesis in compact-object mergers

3 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.

4 r-process Astrophysical sites Core-collapse supernova Neutrino-winds from protoneutron stars. Aspherical explosions, Jets, Magnetorotational Supernova,... [Winteler et al, ApJ 750, L22 (2012); Mösta et al, arxiv: ] Neutron star mergers Matter ejected ( 0.01 M ) dynamically during merger. Electromagnetic emission from radioactive decay of r-process nuclei [KiloNova, Metzger et al (20), Roberts et al (2011), Bauswein et al (2013)] What is the additional contribution from the accretion disk?

5 Role of weak interactions Main processes: ν e + n p + e ν e + p n + e + Neutrino interactions determine the proton to neutron ratio. Neutron-rich ejecta: [ ] E νe E νe > 4 np L νe [ ] L νe 1 E νe 2 np 5 Neutrino cooling and Neutrino-driven wind 4 νe,µ,τ, νe,µ,τ neutron-rich ejecta: r-process R in km 3 Ni proton-rich ejecta: νp-process We need accurate knowledge of ν e and ν e spectra 2 Rns ~ Rν PNS 1.4 n, p νp-process r-process α, p α, p, nuclei α, n α, n, nuclei M(r) in M Si He O 3 νe,µ,τ, νe,µ,τ

6 Neutrino interactions at high densities Most of Equations of State treat neutrons and protons as non-interacting (quasi)particles that move in a mean-field potential U n,p (ρ, T, Y e ). E n = p2 n 2m + m n + U n n µ n µ e E p = p2 p 2m + m p + U p p Q = m n m p + U n U p ν e absorption opacity affected by final state electron blocking ( χ(e ν ) (E ν + m + U) 2 Eν + m ) + U µ e exp, U = U n U p kt ν e absorption affected by energy threshold ( U). χ(e ν ) (E ν m U) 2 µ p E ν > m + U larger symmetry energy (larger U) implies: i) the larger the energy difference between ν e and ν e ; ii) smaller electron flavor luminosities.

7 Constrains in the symmetry energy Combination nuclear physics experiments and astronomical observations (Lattimer & Lim 2013) Isobaric Analog States (Danielewicz & Lee 2013) E sym [MeV] Danielewicz & Lee 2013 IAS Danielewicz & Lee 2013 IAS+Skins Lattimer & Lim 2013 DD2 NL3 TM1 TMA SFHo SFHx FSUgold IUFSU LS180 LS n B [fm -3 ] U [MeV] DD2 NL3 TM1 TMA SFHo SFHx FSUgold IUFSU n B [fm -3 ] Figures from Matthias Hempel (Basel)

8 Impact on neutrino luminosities and Y e evolution 1D Boltzmann transport radiation simulations (artificially induced explosion) for a 11.2 M progenitor based on the DD2 EoS (Stefan Typel and Matthias Hempel). Luminosity [ergs/s] Eν [MeV] Time [s] ν e 11 ν e 9 ν x Time [s] Electron fraction Entropy [kb/baryon] Time [s] Time [s] Y e is moderately neutron-rich at early times and later becomes proton-rich. GMP, Fischer, Huther, J. Phys. G 41, (2014).

9 Nucleosynthesis Rel. abundance Elemental abundance Mass number A Charge number Z Elements between Zn and Mo, including 92 Mo, are produced Mainly neutron-deficient isotopes are produced No elements heavier than Mo (Z = 42) are produced.

10 Neutron decay The neutron-proton energy difference in the medium could be of the order of several s MeV. Neutron decay is an important source of low energy neutrinos. n p + e + ν e e + + n p + ν e This is part of the direct URCA process in neutron stars [Lattimer et al, (1991)] ν e 9 8 Opacity km Ν e e p n Ν e p n e Energy MeV Fischer, Lohs, GMP, Qian, in preparation no neutron decay channel including neutron decay channel t t [s] ν e 20 E ν [MeV] E ν [MeV]

11 Neutron star mergers: Short gamma-ray bursts and r-process Mergers are expected to eject around 0.01 M of very neutron rich-material (Y e 0.01). A similar amount of less neutron-rich material (Y e ) is expected from the accretion disk. They are also promising sources of gravitational waves. Observational signatures of the r-process?

12 Neutron-star mergers: Astrophysically robust ejecta mass fraction Korobkin, Rosswog, Arcones, & Winteler, MNRAS 426, 1940 (2012) Y e [g/cm 3 ] Y e Y i ns1.0-ns1.0 ns1.4-ns1.4 ns1.2-ns1.2 ns1.2-ns1.0 ns1.4-ns1.0 ns1.4-bh5 ns1.4-bh similar results: Bauswein, Goriely, Janka, ApJ 773, 78 (2013) Mass number A

13 General features r-process r-process abundance 200 Proton number (Z) N 1 r, 0 (Si = 6 ) Mass number (A) log(t s 1 ) Neutron number (N) Figure from Peter Möller.

14 Global mass models vs experiment Similar behaviour for all mass models. Problems in reproducing masses in transitional regions.

15 General features evolution in mergers r-process stars once electron fermi energy drops below MeV to allow for beta-decays (ρ 11 g cm 3 ). Important role of nuclear energy production. Increases temperature to values that allow for an (n, γ) (γ, n) equilibrium. r-process operates at moderate high entropies, s 50 0 k/nuc. Trajectories from simmulation A. Bauswein and H.-T. Janka. Log (ρ) (g cm -3 ) T/ 9 K de/dt (Mev nuc -1 s -1 ) FRDM MASS MODEL Time (s)

16 Final abundances different mass models neutron captures computed consistently for each mass model. Final abundances solar r-process abundances R n/seed 1920 FRDM MASS MODEL WS3 MASS MODEL -6 Final abundances HFB21 MASS MODEL DZ31 MASS MODEL A A J. Mendoza-Temis, G. Martinez-Pinedo, K. Langanke, A. Bauswein, H.-Th. Janka, in preparation.

17 Temporal evolution (selected phases) abundances abundances Final abundances Y n /Y h = 1 τ (n, γ) > τ β FRDM MASS MODEL -8-9 solar A Y e Y e abundances abundances Final abundances Y n /Y h = 1 τ (n, γ) > τ β DZ31 MASS MODEL -8-9 solar A Y e Y e

18 Role of N 90 (Telurium isotopes) Sn (MeV) FRDM DZ31 WS3 HFB 21 AME Te (Z=52) Isotopes Neutron Number Hakala et al, PRL 9, (2012) Van Schelt, et al, PRC 85, (2012)

19 Neutron Separation energies Cd isotopes Sn (MeV) FRDM DZ31 WS3 HFB 21 AME 2012 Cd (Z=48) Isotopes Neutron number FRDM mass model predicts rather low neutron separation energies approaching N 90 for Z 50.

20 Odd-even effects (Te isotopes) (3) (MeV) Te (Z = 52) Isotopes FRDM DZ31 WS3 HFB21 AME Neutron Number (3) = ( 1) N [2B e (Z, N) B e (Z, N + 1) B e (Z, N 1)] /2 All mass models have problems reproducing odd-even effects

21 The role of N 130 Sn (MeV) DZ31 WS3 HFB 21 FRDM 1 0 Yb (Z=70) Isotopes Neutron Number Both FRDM and HFB models predict a sudden drop in neutron separation energies approaching N 130 for Z 70.