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Stellar Nucleosynthesis s-Process in Massive Stars
Mounib El Eid American University of Beirut (AUB) Department of Physics
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Physics department American University of Beirut in the desert
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The best gift of the western world to the Middle East
Motto: That they may have life and have it more abundantly. Established: 1866 Type: Private under the law of the state of New York President: Peter F. Dorman Provost: A. Hilal, Provost Staff: 522 full-time instructional faculty Students: 7,036 ( ) Undergraduates: 5,725 Postgraduates: 1,311 Location: Beirut, Lebanon Campus: Urban, 73 acre Website: The best gift of the western world to the Middle East
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outline Introduction 2. Stellar Models s-process results what else?
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We are made from star dust
1. Introductory Comments: Except of the light elements : 1H, 2H, 3He, 4He ,7Be, 7Li All other elements and isotopes were mainly made by nucleosynthesis in stars during their evolution and are expelled during the end stages of the stars. These end stages are: White Dwarfs Neutron Stars Black Holes Stellar nucleosynthesis drives the evolution of stars toward these end stages Stellar nucleosynthesis determines the nuclear energy production in stars and this is what determines their lifetime. Without nuclear energy and gravity: no stars no elements no human beings We are made from star dust
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stars must convert H to He in their interior.
Interesting point: stars must convert H to He in their interior. However, most interstellar helium is of primordial origin. Why? Two reasons: (a) little mixing from the center to the surface of the star, (b) helium is processed into heavier elements during the star’s evolution Red giant Star now Supernova 1987 ongoing cycle Star-forming nebula Supernova Remnant
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A big Q Supernova Remnant E from Radio to X-Ray 40 light years across, 190,000 light years away in the Small Magellanic Cloud
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Expansion at 4 million miles/hour
Kepler's supernova from NASA's Spitzer Space Telescope, Hubble Space Telescope, and e Chandra X-ray observatory Expansion at 4 million miles/hour
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Planetary Nebula The Helix Nebula from La Silla Observatory 700 light years away Will our Sun look like this one day?
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Where are the elements made?
Iron
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IRON Think: iron abundance high due to its high binding energy/nucleon
Mass Number U=cosmological x=cosmic rays (spallations) H=hydrogen burning H=helium burning C=carbon burning O=oxygen burning S=silicon burning E= nuclear statistical equilibrium and r, s, p processes
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Distribution of thee Heavy Elements r, s and p
abundance curve of heavy elements: solar system
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Proton Number Neutron Number Chart of Nuclides; For orientation N =Z
decay + decay Proton Number Proton addition Neutron addition -- decay Neutron Number
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Basics of the s-process
Seed nucleus: Iron Neutron source C(,n)16O AGB stars 22Ne(,n) 25Mg Massive Stars Neutron spectrum Maxwell-Boltzmann Distribution (details later) Temperature range (25 – 30) KeV or (1.5 – 3.5) x 108 K Helium Fusion ( 90 – 100) KeV or ( )x 109 K Carbon Fusion Time scale Beta-decay much faster, Helium fusion Comparable time scale: carbon fusion Massive Stars Thermally pulsating stars (AGB stars Where?
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How does the s-process works, in principle
69Ga 31 s s s s s 30 Z 64Zn 65Zn 66Zn 67Zn 69Zn 70Zn 68Zn s s 29 63Cu 64Cu 65Cu 66Cu s s s s 28 60Ni 61Ni 62Ni 63Ni 64Ni 65Ni s 27 Black: stable (s) 59Co 60Co s s s Red: radioactive (R ) 26 56Fe 57Fe 58Fe 59Fe N
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3.2 Details of the s-process ( Ns – curve)
The s-process works with thermal zed neutrons. They undergo elastic scattering in the star’s plasma. The neutrons obey the Maxwell-Botzmann distribution (MBD): From lab experiments, we know that the neutron-capture cross section varies like: S-wave scattering MBD E0 =most probable energy En E0=kT
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Thermal velocity of the neutron:
Due the the above energy dependence: Let: Maxwell folded Then: Implication: measurement of near vT very useful.
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Time variation of abundances: (Z+1,A)
Radioactive (Z.A) Z N (Z,A-1) A+1 A -1 A with n = neutron flux Beta-rate in stellar plasma (1) Creation destruction
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(Time-integrated neutron flux)
define neutron exposure: Using in Eq. (1) Using: Every radioactive nucleus beta-decays Self-regulating equation with a steady-state solution: This says: a nucleus with small would have small abundance and a nucleus with high would have large abundance.
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Verification of Ns (Si=106 ) (mb) Weak component Massive stars Main component, AGB stars
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Tellur - Isotopes Verification of the With the Te-isotopes 122 123 124 125 126 Te sr s s s sr 121 123 s: s-only r: r-only Sr: both Sb sr sr 124 122 120 Sn sr r r R-process path Advantage of the N-curve: (b) Obtain average neutron flux and temperature ( c) analyze branching (a) To obtain r-abundances: knowing for an isotope , then Ns from the curve and Nr =N-Ns
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Ns (Si=106 ) (mb) Branching
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Stellar Models of Massive Stars
Recent review on this subject: Mounib El Eid, Lih-Sin The & Bradley Meyer: Space Science Review, 2009, online In the folowingI will present : Effect of the 12 C(,) 16 O rate on the evolution including advanced stages (2) Compare different calculations
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Survival of and If 12C(,) would be too effective, no oxygen would survive. But this reaction is very important !!
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The reaction Energy levels above and near the threshold of 12C(,) . For temperatures T9=1.0 and above, the effective stellar energy is near E0=0.30 MeV (Gamow peak). This energy is reached by the low energy tail of the resonance centered at 2.42 MeV (center of mass energy): Two other sub-threshold resonances:
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Several evaluations of
He burning;15-30 MSun CF85: Caughlan et al ; 1985 , At. Data Nucl. Data Tables, 32, 197 Kunz: Kunz etal. : 2002, APJ, 567, 643 NACRE: Angulo et al , Nucl. Phys. A, 656, 3 Buchmann: 1996, APJ, 468, L127
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Details of the stellar models
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Hertzsprung-Russell Diagram
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Energy equation: (Strong, weak, EM interaction
and gravitational. interaction)
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see El Eid, Meyer, The APJ 285, 1994 CONVECTIVE Note that the energy production does not comprise the whole convective core NACRE NACRE The, El Eid, Meyer: APJ, 533 (2007)
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Imbriani et al (2001), ApJ 558, 903 Rate of 25 Msun star CF85 > CF88 CF85 X 12=0.18 No convective carbon-burning core X12 lower Remaining car bon mass fraction No convective core CF88 X12=0.42 But here
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No convective carbon-burning
core Woosley, Heger &Weaver (2002) Rev .Mod. Phys. 74, 1015
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s-process calculations
(a) 22Ne(,n) rate (b) S-process in central helium burning (c ) S-process in C-shell brning
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14N(, ) 18F(e+, ) 18O (, ) 22Ne(,n) 25Mg
Central He-burning NACRE rate smaller than that of CF88 up to T8=2.4 14N(, ) 18F(e+, ) 18O (, ) 22Ne(,n) 25Mg This happens at begin of central helium burning. This becomes effective near end of central He-burning
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s-process in core He-burning
25 M, solar-like initial compositionn Neutron Density Temperature central Helium mass fraction exposure overabundance 25C: CF88 for 25N: NACRE CF88 22Ne(,n) release neutrons earlier in He-burning neutron exposure larger which leads to larger overproduction of 80Kr It does not help to have higher neutron density as in 25N, since this happens late in helium burning
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Results for core He-burning
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(X/X)Kr80 25C 25K 25N R91 pig09 618 183 174 481 232 Pignatari
et al (2009) CF88 Kunz et al NACRE Raiteri et al (91) CF88 NACRE S-process production in this phase not sensitive top 12C(,)16O S-process reduced with the NACRE’s rate The neutron poisoning reaction 16O(n,)17 O with =34 barn at 30 KeV (Igasdhira etal 1995) reduces the 80Kr overabundance by a factor of about 1.8. s-process in core He-burning is less efficient than found before
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S-process in Carbon-Shell burning
The, El Eid, Meyer: APJ, 533 (2007) Pignatari, Gallino, Heil, Wiescher, Kaeppler, Herwig, Bisterzo 2009 in press Raiteri etal (1991),ApJ 367, 228
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NACRE
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Evolution through core Ne-burning phase under different assumptions
Central temperature evolves differently under different assumptions. In 25 K: larger central Ne mass fraction evolution at lower TC but higher C
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S-Process in Carbon shell-burning
Somewhat surprising results and very complicated issue 25 K Kunz et al for C12(,) Peak value
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NACRE rate of C12(,) 25N Peak value 3 episodes
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History of the s-process in the star
Central He-burning Carbon-Shell Helium shell Neutron exposure (not averaged):
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End He-burning C-Shell Burning
Pignatari Ra91a The et al. P/T M2, T9=1.05) (25N) (25K) 23Na 65.8 84.6 53.8 1.22 382 235 315 27Al 2.1 --- 1.16 1.8 140 103 143 37Cl 64.8 -- 72.1 0.9 65.6 61.1 61.8 40K 260 291 268 0.97 137 224 255 50Ti 20 15.9 15.8 1.27 21.6 16.9 16.4 54Cr 14.6 16.5 16.8 0.87 58 Fe 77.2 84.2 105 0.73 55.5 92.8 93.1 59Co 29.9 35.9 36.4 0.82 49.1 39.2 45.5 61Ni 59.9 60.4 67.9 62Ni 53.7 49.9 31.6 1.7 62.5 34.4 36.6 64Ni 166 164 56.6 2.9 221 109 115
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63Cu 134 91.8 58.3 2.3 179 11.5 64.6 65Cu 317 226 122 2.6 224 83.5 148 64Zn 36.8 41 29.4 1.25 4.7 8.6 23.2 66Zn 88.7 118.9 57 1.55 107 62.1 80.6 67Zn 127 171.7 79.4 1.6 272 137 160 68Zn 121 164.7 70 1.72 250 128 131 70Zn 0.3 0.38 0.79 21.7 39.9 70Ge 193 253.7 1.8 421 217 216 72Ge 114 190.7 72.2 1.58 246 187 158 73Ge 128.8 45.1 2.37 381 180 147 74Ge 94.2 99.3 35.9 2.62 251 101 84.9 75AS 60.2 59.6 26.3 2 .29 205 99.1 75.4 76Se 133 212.2 74.7 1.79 333 189 164 80Se 1.9 -- 1.21 95.8 54.7 56.3 80Kr 232 480.7 174 1.33 169 124 181 82Kr 108 210.3 73.4 1.47 86Kr 0.80 ---- 2.57 0.31 73 50.9 10.4 87Rb 1.26 0.63 55.4 43.3 86Sr 147.3 57.1 1.87 103 20.7 138 87Sr 91.4 129.9 47.3 1.93 46.8 28.3 88Sr 26.8 34.8 14.1 39.8 20.1 27.7 96Zr 0.1 --- 0.19 0.53 20.9 11.3 1.18
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Summary of the results for the s-process
Efficiency of the s-process during core He-burning is not sensitive to the 12C(,)16O reaction But sensitive to the 22Ne(,n)25Mg and 16 O(n,)17O reactions (2) Given the present situation, the efficiency of the s-process in massive stars seems to significantly reduced. (3) The s-process in carbon shell-burning depends crucially on the evolution of the star during core helium burning and the following core carbon burning. It is therefore influenced by the xxx 12C(,)16O rate. Interesting: if this rate leads to relatively low mass fraction of carbon at the end of core helium burning , then the s-process occurs later in time, eventually after core Ne-burning. The neutron density achieved is high and the nuclear reaction flow reaches the zirconium region. Otherwise only the Sr region is reached (4) 86 Rb is produced during C-shell burning (5) We estimate: massive stars may contribute at least 40% to the solar–only nuclei with A<87
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In He burning: no time to make errors
In C-shell more time for that
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Concluding Thoughts Nucleosynthesis shows us how the universe is well designed and balanced. We (human being) struggle always to achieve balance in ourselves Stellar nucleosynthesis is driving force of the star’s evolution and a powerful tool to get insight into the internal structure of stars Stellar nucleosynthesis is a basic tool for understanding the chemical evolution of Galaxies Last not least we discover our cosmic connection through the elements that are made in stars. At least in this respect, it is a fascinating branch of Astrophysics. END OF THE LECTURE: THANK YOU for your patience It follows more details about the s-process and some rest material
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Exotic s-process
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Z=2x10-2 : Solar-like Z=10-3 red blue All this happen here
Convective Envelope All this happen here blue red Z=10-3 No Convective Envelope
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Z=10-3 Possible mixing of protons into the
helium shell. It works only at low metallicity. Here: [Fe/H]=-4.5
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Let’s believe it……… then
What is the advantage of this game?
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If this would be true?, we make Primary Sr for example
Result after one time step of proton mixing into the helium convective shell. Strong enhancement near Z=38 57
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LEPP: Light Element Primary Process
These nuclei (Fe-group up to the rising wing of the A=130-peak not reproduced in the supernova wind model (Farouqi et al (2009). They are not explained by the usual secondary ns-process. Farouqi, Kratz, Mashonika, Cowan, Thielemann, Truran (2009)
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Again: Elements with Z<38 are not produced together with those with Z> 38.
The conclusion so far: The light trans-Fe-element are were not produced under the same conditions Eu.
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Stop here
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Stars with different initial metallicity
Observe: For initial stars do not evolve to red giant branch.
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Convective zones of [Fe/H]= 0 25 Ms star
Solar-like composition Convective zones of [Fe/H]= Ms star Z=10 -3 Shaded areas are convective zones. Others are radiative
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Lih-Sin is confusing us: he likes to mix in a most exotic way
Lih-Sin crashed the code Possible mixing of protons into the helium shell. It works only at low metallicity. Here: [Fe/H]=-4.5
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If this would be true?, we make Primary Sr
You have to be more interesting than right Result after one time step of proton mixing into the helium convective shell. Strong enhancement near Z=38 64
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LEPP: Light Element Primary Process
These nuclei (Fe-group up to the rising wing of the A=130-peak not reproduced in the supernova wind model (Farouqi et al (2009). They are not explained by the usual secondary ns-process. Farouqi, Kratz, Mashonika, Cowan, Thielemann, Truran (2009)
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Again: Elements with Z<38 are not produced together with those with Z> 38.
The conclusion so far: The light trans-Fe-element are were not produced under the same conditions Eu.
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r and s processes P s r rs
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Path of R and S processes in Z-N Diagram
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1. Summary of the basic equations of stellar structure& evolution
Dependent variables: u= velocity r=radius T=Temperature =density L=Luminosity P=pressure Independent variables : M r =interior mass t=time Spherical Symmetry is assumed, 1.Momentum equation: (gravitational interaction) 2. Mass conservation 3. Definition of velocity: 4. Energy equation: (Strong, weak & EM interaction ) 5. Energy Transport = mean opacity
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