Astrophysical, observational and nuclear-physics aspects of r-process nucleosynthesis T 1/2 PnPn SnSn nn  Sinaia, 2012 Karl-Ludwig Kratz Part II

Summary “waiting-point” model “weak” r-process “main” r-process (early primary process; SN-II?) superposition of n n -components (later secondary process; explosive shell burning?) seed Fe (implies secondary process) Kratz et al., Ap.J. 662 (2007) …largely site-independent! T 9 and n n constant; instantaneous freezeout To recapitulate… Importance of nuclear-structure input !

Now… from historical, site-independent “waiting-point” approaches to more recent core-collapse SN nucleosynthesis calculations Second part

r-Process calculations with MHD-SN models 2012 … new “hot r-process topic”  magnetohydrodynamic SNe … but, nothing learnt about nuclear-physics input ??? “We investigate the effect of newly measured ß-decay half-lives on r-process nucleosynthesis. We adopt … a magnetohydrodynamic supernova explosion model… The (T 1/2 ) effect slightly alleviates, but does not fully explain, the tendency of r-process models to underpro- duce isotopes with A = 110 – 120…” “We examine magnetohydrodynamically driven SNe as sources of r-process elements in the early Galaxy… … the formation of bipolar jets could naturally provide a site for the strong r-process…” Ap.J. Letter 750 (2012) Phys.Rev. C85 (2012) In both cases FRDM masses have been used partly misleading conclusions

The neutrino-driven wind starts from the surface of the proto-neutron star with a flux of neutrons and protons. As the nucleons cool (≈10 ≥ T 9 ≥ 6), they combine to α-particles + an excess of unbound neutrons. Further cooling (6 ≥ T 9 ≥ 3) leads to the formation of a few Fe-group "seed" nuclei in the so-called α-rich freezeout. Still further cooling (3 ≥ T 9 ≥ 1) leads to neutron captures on this seed compo- sition, making the heavy r-process nuclei. (Woosley & Janka, Nature, 2005) The high-entropy / neutrino-driven wind model Nevertheless, cc-SN “HEW” …still one of the presently favoured scenarios for a rapid neutron-capture nucleosynthesis process

time evolution of temperature, matter density and neutron density extended freezeout phase “best” nuclear-physics input (Mainz, LANL, Basel) Three main parameters: electron abundance Y e = Y p = 1 – Y n radiation entropy S ~ T³/  expansion speed v exp  durations   and  r (Farouqi, PhD Mainz 2005) The Basel – Mainz HEW model nuclear masses β-decay properties n-capture rates fission properties full dynamical network (extension of Basel model)

First results of dynamical r-process calculations Conditions for successful r-process  „strength“ formula Neutron-rich r-process seed beyond N=50 ( 94 Kr, 100 Sr, 95 Rb…)  avoids first bottle-neck in classical model Initial r-process path at particle freezeout (T 9  3)  Y n /Y seed ≤ 150 Total process duration up to Th, U  τ r ≤ 750 ms instead of  3500 ms in classical model Due to improved nuclear-physics input (e.g. N=82 shell quenching!)  max. S  280 to form full 3 rd peak and Th, U Freezeout effects (late non-equilibrium phase)  capture of “free” neutrons  recapture of  dn-neutrons parameters correlated

p α seed Formation of r-process “seed” Time evolution of temperature and density of HEW bubble (V exp =10,000 km/s)  extended “freeze-out” phase! temperature density Recombination of protons and neutrons into α-particles as functions of temperature and time For T 9 ≤ 7  α dominate; at T 9 ≈ 5  p disappear, n survive, “seed“ nuclei emerge. n (Farouqi et al., 2009)

Distribution of seed nuclei effect of different expansion velocities of the hot bubble, V exp distributions are robust ! effect of different electron abundances, Y e distributions show differences, mainly for A > 100 Main seed abundances at A > 80 lie beyond N = 50 !

Parameters HEW model  Y(Z) α n seed Y e =0.45 No neutrons no n-capture r-process! Nucleosynthesis components: S ≤ 100; Y n /Y seed < 1 charged-particle process 100 < S < 150; 1 < Y n /Y seed < 15 “weak” r-process 150 < S < 300; 15 < Y n /Y seed < 150 “main” r-process

Synthesizing the A=130 and A=195 peaks S=196 Process duration: 146 ms S=261 Process duration: 327 ms Pb U Th not yet enough Abundance Y full N r,  distribution  superposition of “weighted” S-components Y e Sτ[ms] Y e Sτ[ms]

Reproduction of N r,  Superposition of S-components with Y e =0.45; weighting according to Y seed No exponential fit to N r,  ! Process duration [ms] Entropy S FRDM ETFSI-Q Remarks A≈115 region top of A≈130 peak REE pygmy peak top of A≈195 peak Th, U fission recycling “ “  significant effect of “shell-quenching” below doubly-magic 132 Sn TTTTTTTT TTTTTTTT

Width of A=130 and 195 peaks ca. 15 m.u.; width of REE pygmy peak ca. 50 m.u.; Widths of all three peaks (A=130, REE and A=195)  S  40 k B /Baryon r-Process robustness due to neutron shell closures Kratz, Farouqi, Ostrowski (2007)

Freezeout at A ≈ 130 Validity and duration of (n,  ) ↔ ( ,n) equilibrium S n (real) / S n (n,  )  1 freezeout Validity and duration of β-flow equilibrium λ eff (Z) x Y(Z) ≠ 0 (n,  -equilibrium valid over  200 ms, independent of Z freezeout Local equilibrium “blobs” with Increasing Z, at different times Definition of different freezeout phases Neutron freezeout at  140 ms, T 9  0.8, Y n /Y r = 1  130 Pd Chemical freezeout at  200 ms, T 9  0.7, Y n /Y r   130 Cd Dynamical freezeout at  450 ms, T 9  0.3, Y n /Y r   130 In

For Y e ≤0.470 full r-process, up to Th, U For Y e  still 3rd peak, but no Th, U For Y e =0.498 still 2nd peak, but no REE Superposition of HEW components ≤ Y e ≤ Farouqi et al. (2009) “weighting” of r-ejecta according to mass predicted by HEW model: for Y e =0.400 ca. 5x10 -4 M  for Y e =0.498 ca M  „What helps…?“ low Y e, high S, high V exp

r-process observables Solar system isotopic N r,  “residuals” T 9 =1.35; n n = r-Process observables today Pb,Bi Observational instrumentation meteoritic and overall solar-system abundances ground- and satellite-based telescopes like Imaging Spectrograph (STIS) at Hubble, HIRES at Keck, and SUBARU recent „Himmelsdurchmusterungen“ HERES and SEGUE Elemental abundances in UMP halo star BD+17°3248 Detection of 33 n-capture elements, the most in any halo star

Observations I: Selected “r-enriched” UMP halo stars Sneden, Cowan & Gallino ARA&A, 2008 Same abundance pattern at the upper end and ??? at the lower end Two options for HEW-fits to UMP halo-star abundances: const. Y e = 0.45, optimize S-range; optimize Y e with corresponding full S-range.

Halo stars vs. HEW-model: Extremes “r-rich” and “r-poor” Eu Factor 25 difference for Sr - Zr region ! gg gg Elemental abundance ratios UMP halo stars r-rich “Cayrel star” / r-rich “Sneden star”r-poor “Honda star” / r-rich “Sneden star” normalized to Eu

r-poor “Honda star” 10 ≤ S ≤ 220 incomplete main r-process 35% N r,  (Eu) r-rich “Sneden star” 140 ≤ S ≤ % N r,  (Eu) full main r-process Extremes “r-rich” and “r-poor”: S-range optimized Sr – Cd region underabundant by a mean factor ≈ 2 relative to SS-r (assumption by Travaglio et al. that this pattern is unique for all UMP halo stars) Sr – Cd region overabundant by a mean factor ≈ 8 relative to SS-r “missing” part to SS-r = LEPP

Halo stars vs. HEW-model: Y/Eu and La/Eu (I. Roederer et al., 2010; K. Farouqi et al., 2010) Instead of restriction to a single Y e with different S-ranges, probably more realistic, choice of different Y e ’s with corresponding full S-ranges 39 Y represents charged-particle component (historical “weak” n-capture r-process) 57 La represents “main” r-process Clear correlation between “r-enrichment” and Y e Caution! La always 100 % scaled solar; log(La/Eu) trend correlated with sub-solar Eu in “r-poor” stars

“… the Ge abundances… track their Fe abundances very well. An explosive process on iron peak nuclei (e.g. the  -rich freezeout in SNe), rather than neutron capture, appears to have been the dominant mechanism for this element…” “Ge abundance seen completely uncorrelated with Eu” Ap.J., 627 (2005) solar abundance ratio Ge – Eu correlation Zr – Eu correlation Correlation between the abundance ratios of [Zr/Fe] (obtained exclusivley with HST STIS) and [Eu/Fe]. The dashed line indicates a direct correlation between Zr and Eu abundances. Correlation between the abundance ratios [Ge/Fe] and [Eu/Fe]. The dashed line indicates a direct correlation between Ge and Eu abundances. As in the previous Figure, the arrow represents the derived upper limit for CS The solid green line at [Ge/Fe]= is a fit to the observed data. A typical error is indicated by the cross. Relative abundances [Ge/H] displayed as a function [Fe/H] metallicity for our sample of 11 Galactic halo stars. The arrow represents the derived upper limit for CS The dashed line indicates the solar abundance ratio of these elements: [Ge/H] = [Fe/H], while the solid green line shows the derived correlation [Ge/H]=[Fe/H]= A typical error is indicated by the cross. Observations: Correlation Ge, Zr with r-process? Can our HEW model explain observations ? “Zr abundances also do not vary cleanly with Eu” Ge okay! 100% CPR Zr two components?

ELEMENT Y(Z) as fct. of S in % 20 ≤ S ≤ ≤ S ≤ ≤ S ≤ Sr Y Zr Pd Ag Te Ba--100 HEW with Y e =0.45; v exp =7500 CP component uncorrelated weak-r component weakly correlated main-r component strongly correlated n-rich α-freezeout normal α -rich freezeout with Eu ? Relative elemental abundances, Y(Z) From Sr via Pd, Ag to Badifferent nucleosynthesis modes

Halo stars vs. HEW model: Sr/Y/Zr as fct. of [Fe/H], [Eu/H] and [Sr/H] Robust Sr/Y/Zr abundance ratios, independent of metallicity, r-enrichment, α-enrichment. Same nucleosynthesis component: CP-process, NOT n-capture r-process Correlation with main r-process (Eu) ? –— HEW model (S ≥ 10); Farouqi (2009) – – average halo stars; Mashonkina (2009)

Zr in r-poor stars overabundant type “Honda star” Zr in halo & r-rich stars underabundant type “Sneden star” Halo, r-rich and r-poor stars are clearly separated! Halo stars vs. HEW model: Zr/Fe/Eu vs. [Eu/Fe], [Fe/H] and [Eu/H] SS Mashonkina (2009); Farouqi (2009) Strong Zr – Eu correlation SS diagonal HEW (av.) r-poor r-rich Zr – Eu correlation Cowan et al. (2005) similar metallicity different r- enrichment

ELEMENT Y(Z) as fct. of S in % 20 ≤ S ≤ ≤ S ≤ ≤ S ≤ Sr Y Zr Pd Ag Te Ba--100 HEW with Y e =0.45; v exp =7500 CP component uncorrelated weak-r component weakly correlated main-r component strongly correlated n-rich α-freezeout normal α -rich freezeout with Eu ? Relative elemental abundances, Y(Z) From Sr via Pd, Ag to Badifferent nucleosynthesis modes

Halo stars vs. HEW-model predictions: Pd average Halo stars -1.5 < [Eu/H] < -0.6 log(Pd/Sr)  (0.09) HEW model log(Pd/Sr) = (“r-rich”) (“r-poor”) average Halo stars -1.5 < [Eu/H] < -0.6 log(Pd/Eu)  +0.81(0.07) HEW model log(Pd/Eu) = (Pd CP+r) (“r-rich”) (”r-poor”) r-poor stars ([Eu/H] < -3) indicate TWO nucleosynthesis components for 46 Pd: Pd/Sr  uncorrelated, Pd/Eu  (weakly) correlated with “main” r-process Pd relation to CP-element Sr Pd relation to r-element Eu Pd/Sr and Pd/Eu different from Sr/Y/Zr

Observations of Pd & Ag in giant and dwarf stars C. Hansen et al.; A&A in print (2012) indication of different production processes (Sr – charged-particle, Ag – weak-r, Eu – main r-process) anticorrelation between Ag and Sranticorrelation between Ag and Eu correlation of Pd and Ag Pd and Ag are produced in the same process (predominantly) weak r-process All observations in agreement with our HEW predictions !

CS31082 CS22892 Observations: [Ag/Fe] vs. [Eu/Fe] BD+17°3248 HD HD HD HD88609 HD “r-poor” “r-rich” Selection of pure r-process stars; no measurable s-process “contamination”

SAGA data base: [X/Fe] vs. [Eu/Fe] (I) Zn Sr Zr SS r-poor r-rich Transition from CP-component to “weak” n-capture r-process Ag

Th SAGA data base: [X/Fe] vs. [Eu/Fe] (II) m1m1 Ir m1m1 Dy m1m1 SS …Th – U – Pb cosmochronology ?!

The SS A ≈ 130 r-peak r-Process observables today Observational instrumentation meteoritic and overall solar-system abundances ground- and satellite-based telescopes like Imaging Spectrograph (STIS) at Hubble, HIRES at Keck, and SUBARU recent „Himmelsdurchmusterungen“ HERES and SEGUE From elemental abundances to isotopic yields …; More sensitive to nuclear physics input, here in the 120 < A < 140 region Xe-H shows strong depletion of lighter isotopes with s-only 128, 130 Xe “zero”. r-process observables

Xe in nanodiamond stardust (I) (U. Ott)

Xe in nanodiamond stardust (II) (U. Ott)

Xe in nanodiamond stardust (III)

The “neutron-burst“ model (I) The “neutron-burst” model is the favored nucleosynthesis scenario in the cosmochemistry community; so far applied to isotopic abundances of Zr, Mo, Ru, Te, Xe, Ba, Pt Historical basis: …neutron burst in shocked He-rich matter in exploding massive stars first ideas by Cameron Howard et al.; Meteoritics 27 (1992) Rauscher et al.; Ap.J. 576 (2002) Several steps: 1)start with SOLAR isotope distribution 2)exposure to weak neutron fluence (  = 0.02 mbarn -1 ) mimics weak s-processing during pre-SN phase 3)s-ashes heated suddenly to T 9 = 1.0 by SN shock 4)expansion and cooling on 10 s hydrodynamical timescale resulting n-density (from ( ,n)-reactions) during burst ≈ n/cm 3 for ≈ 1 s final n-exposure mbarn -1 shift of post-processed isotope pattern towards higher A strong time and temperature dependence

Howard‘s Xe-H abundances N* are pure, unmixed yields; in addition, arbitrary mix with 5.5 parts N סּ => N-H N-H = (N* + 5.5N סּ )/5.5N סּ = N*/N סּ N-H = 1 s-only 128,130 Xe totally eliminated …however, “there is little physical reason to admix N* with solar abun- dances, other than the tradition of interpreting isotopic anomalies as a binary admixture with SS abundances.“ arbitrary N סּ -mixing somewhat too strong! The “neutron-burst“ model (II)

Two SN models a) site-independent “waiting-point“ approach b) high-entropy-wind ejecta of core-collapse SN-II full parameter space in a) n n -components b) S-ranges No reproduction of Xe-H ! special variants of r-process? “hot“ “cold“ r-process “strong” Xe in SN-II

The A  130 SS r-abundance peak HEW reproduction of the solar-system A  130 r-abundance peak Nuclear-physics input: ETFSI-Q; updated QRPA(GT+ff); experimental data Y e = 0.45 “cold” r-process V exp = 20000; S(top)  140 peak top at A  126 “hybrid” r-process V exp = 7500; S(top)  195 peak top at 128 < A < 130 “hot” r-process V exp = 1000; S(top)  350 peak top at A  136 as, e.g. in Arcones & Martinez-Pinedo, Phys. Rev. C83 (2011)

Successive “cuts“ of low-S components exclusion of “weak“ r-process “strong“ r-process ! Xe-H – HEW “old“ “old“ nuclear physics input e.g. 129 Ag T 1/2 (g.s.) = 46 ms; 130 Pd T 1/2 = 98 ms; P 1-2n = 96 %; 4 %

Experiment: T 1/2 = 68 ms; P n = 3.4 % Surprising  -decay properties of 131 Cd QRPA predictions: T 1/2 (GT) = 943 ms; P n (GT) = 99 % T 1/2 (GT+ff) = 59 ms; P n (GT+ff) = 14 % Remember: …just ONE neutron outside N=82 magic shell Nuclear-structure requests: higher Q β, lower main-GT, low-lying ff-strength “new” nuclear-physics input for A ≈ 130 peak region

…but, still „cuts“ of low-S region necessary again, exclusion of “weak“ r-process “strong“ r-process ! 129 Xe, 131 Xe, 134 Xe okay; 132 Xe still factor ≈ 2 too high Xe-H – HEW “new“ “new“ nuclear physics input optimized to measured T 1/2, P n and decay schemes of 130 Cd and 131 Cd e.g. 129 Ag stellar-T 1/2 = 82 ms; 130 Pd T 1/2 (new/old) ≈ 0.2; P 1-2n (new/old) ≈ 0.8; 4.5 reduction of 129 Xe increase of 136 Xe

…same HEW r-process conditions as Xe-H ! towards Rauscher & N ʘ L'09 towards HEW Pt-H in nanodiamonds … at the end, an easy case for HEW SN-II “ cold “ r-process ! recent AMS measurement by Wallner et al.

Summary Still today there is no selfconsistent hydro-model for SNe, that provides the necessary astrophysical conditions for a full r-process parameterized dynamical studies (like our HEW approach) are still useful to explain r-process observables; astronomical observations & HEW calculations indicate that SS-r and UMP halo-star abundance distributions are superpositions of 3 nucleosynthesis components:charged-particle, weak-r and main-r the yields of the CP-component (up to Zr) are largely uncorrelated with the “main” r-process; the yields of the weak-r component (Mo to Cd) are partly correlated with the “main” r-process; elements ≥ Te belong to the “main” r-process HEW can explain the peculiar isotopic anomalies in SiC-X grains and nanodiamond stardust Therefore,

K. Farouqi (Mainz) B. Pfeiffer (Mainz & Gießen) O. Hallmann (Mainz) U. Ott (Mainz) P. Möller (Los Alamos) F.-K. Thielemann (Basel) W.B. Walters (Maryland) L.I. Mashonkina (RAS, Moscow) C. Sneden & I. Roederer (Austin, Texas) A. Frebel (MIT, Cambridge) N. Christlieb (LSW, Heidelberg) C.J. Hansen (LSW, Heidelberg) Main collaborators:

Discussion

??? Neutron star mergers

P n effects at the A≈130 peak Significant differences (1) smoothing of odd-even Y(A) staggering (2) importance of individual waiting-point nuclides, e.g. 127 Rh, 130 Pd, 133 Ag, 136 Cd (3) shift left wing of peak in clear contrast to Arcones & Martinez-Pinedo Phys. Rev. C83 (2011)

Halo stars vs. HEW model: “LEPP” elements HEW (10 < S < 280) WP (Fe seed; < n n < ) weak-r (Si seed; n n  ) LEPP-abundances vs. CP- enrichment (Zr) HEW reproduces high-Z LEPP observations (Sr – Sn); underestimates low-Z LEPP observations (Cu – Ge) additional nucleosynthesis processes ? (e.g. Fröhlich et al. p-process; Pignatari et al. rs-process; El Eid et al. early s-process) (Farouqi, Mashonkina et al. 2008)

Th/U – a robust and reliable r-process chronometer? Neutron-density range [n/cm 3 ] r-Chronometer age [Gyr] norm ± 2.8 Th/U robust – yes reliable – no (not always) Correlate Th, U with other elements, e.g. 3rd peak, Pb !

How about Pb as r-chronometer ? In principle, direct correlation of Pb to Th, U ≈ 85% abundance from α-decay from actinide region Sensitivity to age: Pb 0 Gyr Y(Pb) = (Mainz) // Sneden et al. (2008) Y(Pb) = ? 5 Gyr Gyr Th,U/Pb 13 Gyr: Th/Pb ≈ U/Pb ≈ New HEW results: 1)agreement Th/U and Th/Pb with NLTE analysis of Frebel-star (HE ; [Fe/H] ≈ -3) 2)consistent picture for Th/Ba – Th/U for 4 “r-rich” halo-stars with Y e ≈ 0.45; and for the “actinide-boost” Cayrel-star with Y e ≈ 0.44.

Effect of n-capture rates on A=130 and A=195 peaks blue curve: green curve: 1/100 x red curve: 100 x at A=130 peak fast freezeout  no significant effect at A=195 peak slow freezeout  significant effect T-dependent NON-SMOKER rates; ETFSI-Q + AME2003