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The origin of heavy elements in the solar system

The origin of heavy elements in the solar system. (Pagel, Fig 6.8). each process contribution is a mix of many events !. Heavy elements in Metal Poor Halo Stars. CS22892-052 red (K) giant located in halo distance: 4.7 kpc mass ~0.8 M_sol [Fe/H]= - 3.0 [Dy/Fe]= +1.7.

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The origin of heavy elements in the solar system

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  1. The origin of heavy elements in the solar system (Pagel, Fig 6.8) each process contribution is a mix of many events !

  2. Heavy elements in Metal Poor Halo Stars CS22892-052red (K) giantlocated in halodistance: 4.7 kpcmass ~0.8 M_sol [Fe/H]= -3.0 [Dy/Fe]= +1.7 recall:[X/Y]=log(X/Y)-log(X/Y)solar old stars - formed before Galaxy was mixedthey preserve local pollution from individual nucleosynthesis events

  3. A single (or a few) r-process event(s) CS22892-052 (Sneden et al. 2003) solar r CosmoChronometer abundance log(X/H)-12 NEW:CS31082-001 with U(Cayrel et al. 2001) + Age: 16 3 Gyr(Schatz et al. 2002 ApJ 579, 626) - element number other, secondr-process to fillthis up ?(weak r-process) main r-processmatches exactly solar r-pattern conclusions ?

  4. Overview heavy element nucleosynthesis

  5. neutron capture timescale: ~ 0.2 ms Rapid neutroncapture b-decay Equilibrium favors“waiting point” (g,n) photodisintegration The r-process Temperature: ~1-2 GK Density: 300 g/cm3 (~60% neutrons !) Seed Proton number Neutron number

  6. show movie

  7. Waiting point approximation Definition: ASSUME (n,g)-(g,n) equilibrium within isotopic chain How good is the approximation ? This is a valid assumption during most of the r-process BUT: freezeout is neglected Freiburghaus et al. ApJ 516 (2999) 381 showed agreement with dynamical models Consequences During (n,g)-(g,n) equilibrium abundances within an isotopic chain are given by: • time independent • can treat whole chain as a single nucleus in network • only slow beta decays need to be calculated dynamically • neutron capture rate independent(therefore: during most of the r-process n-capture rates do not matter !)

  8. Endpoint of the r-process r-process endedby n-induced fission or spontaneousfission (different pathsfor different conditions) (Goriely & Clerbaux A&A 348 (1999), 798 b-delayed fission spontaneous fission n-induced fission n-capture (DC) b- (Z,A) (Z,A) fission fission fission (Z,A+1) fission barrier (Z,A+1) (Z+1,A)

  9. Consequences of fission Fission produces A~Aend/2 ~ 125 nuclei modification of abundances around A=130 peak fission products can serve as seed for the r-process - are processed again into A~250 region via r-process - fission again fission cycling ! Note: the exact endpoint of the r-process and the degree and impact of fission are unknown because: • Site conditions not known – is n/seed ratio large enough to reach fission ? (or even large enough for fission cycling ?) • Fission barriers highly uncertain • Fission fragment distributions not reliably calculated so far (for fission from excited states !)

  10. Role of beta delayed neutron emission Neutron rich nuclei can emit one or more neutrons during b-decay if Sn<Qb (the more neutron rich, the lower Sn and the higher Qb) b- (Z,A) n Sn (Z+1,A-1) g (Z+1,A) If some fraction of decay goes above Sn in daughter nucleusthen some fraction Pnof the decays will emit a neutron (in addition to e-and n)(generally, neutron emission competes favorably with g-decay - strong interaction !)

  11. during r-process: none as neutrons get recaptured quicklyduring freezeout : Effects: • modification of final abundance • late time neutron production (those get recaptured) Calculated r-process production of elements (Kratz et al. ApJ 403 (1993) 216): after b-decay before b-decay smoothing effect from b-delayed n emission !

  12. Example: impact of Pn for 137Sb A=136 ( 99%) A=137 ( 0%) Cs (55) r-processwaiting point Xe (54) Pn=0% I (53) Te (52) Pn=99.9% Sb (51) Sn (50) In (49) Cd (48) Ag (47) r-process waiting point

  13. Summary: Nuclear physics in the r-process

  14. First NSCL experiment The r-process path r-process abundance distribution r-processpath RIA Reach New MSU/NSCL Reach Known (Reach for half-life)

  15. National Superconducting Cyclotron Laboratory atMichigan State University New Coupled Cyclotron Facility – experiments since mid 2001 Ion Source:86Kr beam 86Kr beam140 MeV/u Tracking (=Momentum) TOF start Implant beam in detectorand observe decay 86Kr hits Be target and fragments TOF stop dE detector Separated beamof r-processnuclei Fast beam fragmentation facility – allows event by event particle identification

  16. Installation of D4 steel, Jul/2000

  17. First r-process experiments at new NSCL CCF facility (June 02) • Measure: • b-decay half-lives • Branchings for b-delayed n-emission • Detect: • Particle type (TOF, dE, p) • Implantation time and location • b-emission time and location • neutron-b coincidences New NSCL Neutron detectorNERO 3He + n -> t + p neutron Fast Fragment Beam Si Stack (fragment. 140 MeV/u 86Kr)

  18. NSCL BCS – Beta Counting System • 4 cm x 4 cm active area • 1 mm thick • 40-strip pitch in x and y dimensions ->1600 pixels Si Si Si BCS b

  19. NERO – Neutron Emission Ratio Observer BF3 Proportional Counters 3He Proportional Counters • Specifications: • 60 counters total • (16 3He , 44 BF3) • 60 cm x 60 cm x 80 cm • polyethylene block • Extensive exterior • shielding • 43% total neutron • efficiency (MCNP) Polyethylene Moderator Boron Carbide Shielding

  20. June 2002 Data – preliminary results Mainz: K.-L. Kratz, B. Pfeiffer PNNL: P. Reeder Maryland/ANL: W.B. Walters, A. Woehr Notre Dame: J. Goerres, M. Wiescher NSCL:P. Hosmer, R. Clement, A. Estrade, P.F. Mantica, F. Montes, C. Morton, M. Ouellette, P. Santi, A. Stolz r-processpath 77Ni 74Co 71Fe Gated b-decay time curve(implant-decay time differences) Energy loss Time of flight Fast RIBs: • cocktail beams • no inflight decay losses • measure with low rates (>1/day)

  21. Neutron Data With neutron in addition Nuclei with decay detected 420 420 Nn 370 370 DE (arb units) DE(arb units) 320 320 76Ni 76Ni 73Co 73Co 270 270 220 220 450 500 550 350 400 450 500 550 350 400 TOF (arb units) TOF (arb units) neutron detection efficiency (neutrons seen/neutrons emitted)

  22. Overview of common r process models • Site independent models: • nn, T, t parametrization (neutron density, temperature, irradiation time) • S, Ye, t parametrization (Entropy, electron fraction, expansion timescale) • Core collapse supernovae • Neutrino wind • Jets • Explosive helium burning • Neutron star mergers

  23. Site independent approach Goal: Use abundance observations as general constraints on r-process conditions BUT: need nuclear physics to do it nn, T, t parametrization (see Prof. K.-L. Kratz transparencies) obtain r-process conditionsneeded for which the right N=50 andN=82 isotopes are waiting points(A~80 and 130 respectively) often in waiting point approximation Kratz et al. ApJ403(1993)216

  24. S, Ye, t parametrization • Consider a blob of matter with entropy S, electron abundance Ye in NSE • Expand adiabatically with expansion timescale t • Calculate abundances - what will happen: • NSE • QSE (2 clusters: p,n,a and heavy nuclei) • a-rich freezeout (for higher S) (3a and aan reactions slowly move matter from p,n,a clusterto heavier nuclei – once a heavy nucleus is created it rapidly captures a-particlesas a result large amounts of A~90-100 nuclei are producewhich serve as seed for the r-process • r-process phaseinitially: n,g – g,n equilibriumlater: freezeout

  25. cross : most abundant nucleuscolors: degree of equilibrium with that nucleus (difference in chemical potential) Evolution of equilibria: from Brad Meyers website

  26. 2 possible scenarios: 1) high S, moderate Ye2) low S, low Ye Results forneutron to seed ratios:(Meyer & Brown ApJS112(1997)199) • n/seed is higher for • lower Ye(more neutrons) • higher entropy(more light particles, less heavy nuclei – less seeds) (or: low density – low 3a rate – slow seed assembly) • faster expansion(less time to assemble seeds)

  27. r-process in Supernovae ? Most favored scenario for high entropy: Neutrino heated wind evaporating from proto neutron star in core collapse ne neutrino sphere (ne+p  n+e+ weak opacity because only few protons present) ne neutrino sphere (ve+n  p+e+ strong opacity because many neutrons present) protoneutron star(n-rich) weak interactions regulate n/p ratio: faster as necome from deeperand are therefore hotter ! ne+p  n+e+ ne+n  p+e- therefore matter is drivenneutron rich

  28. Results for Supernova r-process A~90 overproduction can’t produce A~195 anymore density artificially reduced by factor 5 density artificially reduced by factor 5.5 artificial parameter to getA~195 peak (need S increase) other problem: the a effect Takahashi, Witti, & Janka A&A 286(1994)857(for latest treatment of this scenario see Thompson, Burrows, Meyer ApJ 562 (2001) 887)

  29. r-process in neutron star mergers ?

  30. Ejection of matter in NS-mergers Rosswog et al. A&A 341 (1999) 499 Destiny of Matter:red: ejected blue: tailsgreen: disk black: black hole (here, neutron stars areco-rotating – tidally locked)

  31. r-process in NS-mergers large neutron/seed ratios, fission cycling ! But: Ye free parameter …

  32. Summary theoretical scenarios both seem to be viable scenarios (but discussion ongoing)

  33. r- and s-process elements in stars with varying metallicity • s-process: • later (lower mass stars) • gradual onset (range of stars) • r-process: • very early (massive stars) • sudden onset (no low mass star contribution) confirms massive starsas r-process sites(but includes SN and NS-mergers) (Burris et al. ApJ 544 (2000) 302) s-process r-process ~age

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