Tariq Al-Abdullah Hashemite University, Jordan Cairo 2009

1. Problems and Issues in Nuclear Astrophysics Tariq Al-Abdullah Hashemite University, Jordan Cairo 2009

2. SUMMARY • Why nuclear physics in astrophysics? • Why indirect measurements of cross sections in nuclear astrophysics? • Theindirect teqhniques and their applicatoins! • Perspectives of the method: RIB and (part,g) reaction • Nuclear Physics Research and Education!

3. Questions: When did it start ?? What elements are produced and can we understand the isotopic composition ?? Do parameters of the Early Universe have an influence ?? Nucleosynthesis 1945: Gamow’s Hypothesis “ all of today elements were made during the early BIG BANG of the Universe “ wrong!

4. Big BangNucleosynthesis Wrong for 3 simple reasons • Binding energy of deuteron (2.22 MeV) is too small !! • Binding energy of 4He is too large (28.3 MeV) !! • There are no stable isotopes with A=5 and A=8 !! Deuterons are being dissociated until the Universe Has cooled down to 80 keV !!! for further fusion the train has Long left the station !! Universe composition: ~76% H and ~23% 4He

6. Stellar life cycle BIRTH gravitational contraction Interstellargas Stars mixing of interstellar gas thermonuclear reactions • energy production • stability against collapse • synthesis of “metals” abundance distribution explosion DEATH

7. Experimental nuclear astrophysics The goals of experimental Nuclear Astrophysics are: - study of the origin of the elements or nucleosynthesis • - studyof the energy generation processes in stars. M. Smith & E. Rehm Courtesy: M. Arnould

8. Thermonuclear reactions in stars: general features Reaction rate:r = N1 N2 v (v)(# reactions Volume-1 Time-1) In stellar plasma the velocity of particles varies over a wide range Assume (v) is the velocity distribution  Average Reaction rate per particle pair:

9. Thermonuclear reactions in stars: general features <v> is a KEY quantity Total reaction rate R12 = (1+12)-1 N1N2<v>12 Energy production rate 12 = R12 Q12 Mean lifetime of nuclei X against destruction by nuclei a energy productionas star evolves - change in abundanceof nuclei X <sv> to be determined from experiments / theoretical considerations as star evolves, T changes evaluate <sv> for each temperature NEED ANAYLITICAL EXPRESSION FOR !

10. Thermonuclear reactions in stars:charged particles charged particles  Coulomb barrier V energy available from thermal motion Coulomb potential ECoul ~ Z1Z2 (MeV) Ekin ~ kT (keV) tunnel effect r0 r nuclear well T ~ 15x106 K (e.g. our Sun)  kT ~ 1 keV T ~ 1010 K (Big Bang) kT ~ 1 MeV reactions occur by TUNNEL EFFECT tunneling probabilityP  exp(-2) during quiescent burnings: kT << ECoul exp(-2ph) = GAMOW factor in numerical units:2ph = 31.29 Z1Z2(m/E)½(m in amu and Ecm in keV)

11. Thermonuclear reactions in stars:Astrophysical factor For non-resonant reactions, the cross section behaviour is dominated by the Gamow factor Sharp drop with energy!!!! Cross section can be parameterized as (E)=(1/E)exp(-2)S(E) De Broglie wavelenght Penetration probability ASTROPHYSICAL FACTOR S(E) is a sort of linearization of the cross section where all non-nuclear effects have been taken out

12. (E)  exp(-E/kT) Probability (E) (E)  E kT Energy Thermonuclear reactions in stars: general features Quiescent stellar burning scenarios:non-relativistic, non-degenerate gasin thermodynamic equilibrium at temperature T Maxwell-Boltzmann velocitydistribution  = reduced mass v = relative velocity Reaction rate:

13. Thermonuclear reactions in stars: Gamow window varies smoothly with energy Gives the energy dependence MAXIMUM reaction rate: tunnelling through Coulomb barrier  exp(-) Maxwell-Boltzmann distribution  exp(-E/kT) E0 < E0 Gamow peak relative probability only small energy range contributes to reaction rate  OK to set S(E) ~ S(E0) = const. E0 energy kT E0

14. Experimental approach: generalities Scenario of quiescent burning stages of stellar evolution FEATURES • T ~ 106–108 KE0 ~ 100 keV<<Ecoul  tunnel effect • 10-20 barn <  < 10-9 barn •  average interaction time  ~ <v>-1 ~ 109 y •  unstable species DO NOT play a significant role PROBLEMS 10-20 b <  < 10-9 b poor signal-to-noise ratio  major experimental challenge extrapolation procedure required REQUIREMENTS poor signal-to-noise ratio long measurements  ultra pure targets  high beam intensities  high detection efficiency

15. Experimental approach: extrapolation Experimental procedure • measure (E) over a wide range of energies, • EXTRAPOLATE down to Gamow energy region around E0 (E) resonance LOG SCALE non-resonant many orders of magnitude direct measurements C.M. Energy E0 Ecoul extrapolation needed ! Coulomb barrier

16. Experimental approach: extrapolation II Even using the S(E)-factor,extrapolation is not a piece of cake!!! S(E) directmeasurement (LINEAR SCALE) low-energy tail of broad resonance Extrapolation non resonant process sub-threshold resonance interaction energy E Er -Er 0 DANGER IN EXTRAPOLATION:large uncertainties!

17. Experimental approach: avoiding extrapolation IDEA: To avoid extrapolationit is necessary to measure Cros sections in theGamow region EXPERIMENTAL SOLUTION • IMPROVEMENTS TO INCREASE THE NUMBER OF DETECTED PARTICLES • New accelerator with high beam intensity • Gas target • 4 pdetectors IMPROVEMENTS TO REDUCE THE BACKGROUND • Use of laboratory with natural shield reduce (cosmic) background example:LUNAfacilityin Italy

18. Experimental approach: new problem BUT… at astrophysical energies A NEW PROBLEM ARISE The ELECTRON SCREENING WHY IS THIS A PROBLEM? It is a problem because electron screening in STARS and in LABORATORIES is not the same!

19. To extract the bare astrophysical Sb(E)–factor fromdirect (shielded) measurements extrapolation were performed at higher energies To avoid extrapolation experimental techniques were improved to perform measurement at very low energies After improving measurements at very low energies, electron screening effects were discovered EXTRAPOLATION IS BACK AGAIN Is there any way out ?

20. Coulomb Dissociationmethod(radiative capture reactions). INDIRECT METHODS In order to solve some of the problem cited above (low cross sections, electron screening) some indirect approaches were proposed such as: Asymptotic Normalisation Coefficients (ANC) method (radiative capture reactions). Trojan Horse Method(thermonuclear reactions induced by light particles)

21. A B (A+p) Y X (Y+P) Asymptotic Normalization Coefficients (ANCS) • Direct Capture Reactions for charges particles: • The binding energy of the captured particle is low. • The capture occurs through the tail of the overlap function. • The Amplitude of the tail is given by the ANCs. • For a Transfer reaction (X+A→Y+B): • The DWBA amplitude: • The Asymptotic behavior of the radial overlap function: • The Asymptotic normalization of the bound-state wave function: • For r > RN, the radial dependences are the same

22. A C2(B) B(A+a) a Y X (Y+a) C2(X) Extracting the ANCS • Peripheral Transfer Reaction (X+A→Y+B): • The reaction cross section: • In terms of the ANCs: • Procedure to extract the ANCs:

23. Experiments using the ANCS 23Al 24Si 24Al 25Si 26Si 25Al 20Mg 21Mg 22Mg 23Mg 24Mg (p,γ) (p,α) (β+ ν) = studied at TAMU Comp with direct meas: 16O(3He,d)17F vs. 16O(p,g)17F Gagliardi e.a. PRC 1999 vs. Morlock e.a. PRL 1997 June 2008 19Na 20Na 23Na 21Na 22Na 17Ne 18Ne 19Ne 20Ne 21Ne 22Ne 17F 18F 19F 16F 15F 15O 16O 17O 18O 13O 14O 12N 13N 14N 15N 11N 11C 12C 13C 10C 9C 7Be(p,g)8B (solar neutrinos probl.): p-transfer:S17(0)=18.2±1.7 eVb Breakup: S17(0)=18.7±1.9 eVb Direct meas: S17(0)=20.8±1.4 eVb 8B 9B 10B 11B CNO, HCNO 8Be 9Be 7Be Ne-Na cycle

24. ANC’s measured by stable beams • 9Be + p« 10B [9Be(3He,d)10B;9Be(10B,9Be)10B] • 7Li + n« 8Li [12C(7Li,8Li)13C] • 13C+p«14N [13C(3He,d)14N;13C(14N,13C)14N] • 14N + p« 15O [14N(3He,d)15O] • 16O + p« 17F [16O(3He,d)17F] • 20Ne+p«21Na [20Ne(3He,d)21Na] • beams»10 MeV/u

25. ANC’s measured by radioactive (rare isotope)beams • 7Be + p« 8B [10B(7Be,8B)9Be] • [14N(7Be,8B)13C] • 11C+p«12N [14N(11C,12N)13C] • 13N+p«14O [14N(13N,14O)13C] • 17F + p« 18Ne [14N(17F,18Ne)13C] • beams»10 - 12 MeV/u

26. ANC’s measured by stable beams (mirror symmetry) • 7Be+p«8B [13C(7Li,8Li)12C] • 22Mg+p«23Al [13C(22Ne,23Ne)12C]** • 17F+p«18Ne [13C(17O,18O)12C]** ** T. Al-Abdullah, PhD Thesis

27. Rare Isotope Accelerators

28. Why RIA ?? • How are the heavy elements created? • How do nuclear properties influence the stars? • What is the structure of atomic nuclei? • How do complex systems get properties from their constituents? • How can complex many-body systems display regularities? • Which new symmetries characterize exotic nuclei? • What are the fundamental symmetries of nature?

29. Radioactive Nuclei in Supernovae

35. International Prespectives • The international effort to study the science of rare isotopes is highly complementary. • RIA will be the first and only facility that will have the capability to meet the challenge of understanding the origin of the elements. • RIA will attract the brightest minds, new generations of the highest-caliber students and the future nuclear scientists. • RIA will provide many new isotopes that can be used to specific diagnostic and therapeutic applications. RIA: Connecting Nuclei with the Universe

36. Research and Education

37. Thank you