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

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

Nucleosynthesis 1945 : Gamow’s Hypothesis “ all of today elements were made during the early BIG BANG of the Universe “ 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 ?? wrong!

Big Bang Nucleosynthesis Wrong for 3 simple reasons oBinding energy of deuteron (2.22 MeV) is too small !! oBinding energy of 4 He is too large (28.3 MeV) !! oThere 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% 4 He

Stellar life cycle energy production  energy production  stability against collapse  synthesis of “metals ” thermonuclearreactions BIRTH gravitational contraction explosion DEATH mixing of interstellar gas Interstellar gasStars abundance distribution

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

 Average Reaction rate per particle pair: In stellar plasma the velocity of particles varies over a wide range Assume  (v) is the velocity distribution Thermonuclear reactions in stars: general features Reaction rate: r = N 1 N 2 v  (v) ( # reactions Volume -1 Time -1 )

Thermonuclear reactions in stars: general features is a KEY quantity Total reaction rate R 12 = (1+  12 ) -1 N 1 N 2 12 Energy production rate  12 = R 12 Q 12 to be determined from experiments / theoretical considerations as star evolves, T changes  evaluate for each temperature energy production as star evolves - change in abundance of nuclei X NEED ANAYLITICAL EXPRESSION FOR  ! Mean lifetime of nuclei X against destruction by nuclei a

charged particles  Coulomb barrier tunnel effect E kin ~ kT (keV) E Coul ~ Z 1 Z 2 (MeV) nuclear well Coulomb potential V rr0r0 T ~ 15x10 6 K (e.g. our Sun)  kT ~ 1 keV T ~ K (Big Bang)  kT ~ 1 MeV energy available from thermal motion exp(-2  = GAMOW factor reactions occur by TUNNEL EFFECT tunneling probability P  exp(-2  ) during quiescent burnings: kT << E Coul Thermonuclear reactions in stars: charged particles in numerical units: 2  = Z 1 Z 2 (  /E) ½  i  n amu and E cm in keV)

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

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

Thermonuclear reactions in stars: Gamow window E0E0 Gives the energy dependence MAXIMUM reaction rate: Gamow peak tunnelling through Coulomb barrier  exp(- ) Maxwell-Boltzmann distribution  exp(-E/kT) relative probability energy kTE0E0 varies smoothly with energy only small energy range contributes to reaction rate  OK to set S(E) ~ S(E 0 ) = const.  E 0 < E 0

T ~ 10 6 – 10 8 K  E 0 ~ 100 keV << E coul  tunnel effect  barn <  < barn  average interaction time  ~ -1 ~ 10 9 y  unstable species DO NOT play a significant role Scenario of quiescent burning stages of stellar evolution FEATURES PROBLEMS b <  < 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 Experimental approach: generalities

 measure  (E) over a wide range of energies,  EXTRAPOLATE down to Gamow energy region around E 0 Experimental procedure LOG SCALE E0E0 E coul Coulomb barrier  (E) non-resonant resonance direct measurements extrapolation needed ! many orders of magnitude C.M. Energy Experimental approach: extrapolation

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

EXPERIMENTAL SOLUTION IMPROVEMENTS TO INCREASE THE NUMBER OF DETECTED PARTICLES  New accelerator with high beam intensity  Gas target  4  detectors IMPROVEMENTS TO REDUCE THE BACKGROUND  Use of laboratory with natural shield  reduce (cosmic) background example: LUNA facility in Italy Experimental approach: avoiding extrapolation IDEA: To avoid extrapolation it is necessary to measure Cros sections in the Gamow region

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

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 To extract the bare astrophysical S b (E) –factor from direct (shielded) measurements extrapolation were performed at higher energies extrapolation were performed at higher energies EXTRAPOLATION IS BACK AGAIN Is there any way out ?

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

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 > R N, the radial dependences are the same A Y B (A+p) X (Y+P) Asymptotic Normalization Coefficients (ANCS)

Peripheral Transfer Reaction (X+A → Y+B):  The reaction cross section:  In terms of the ANCs: Procedure to extract the ANCs: A Y B(A+a) X (Y+a) a C 2 (B) C 2 (X) Extracting the ANCS

Ne-Na cycle 12 C 13 C 13 N 15 N 15 O 14 N 17 O 17 F 16 O 19 F 18 F 18 O 14 O 19 Ne 18 Ne 13 O 11 C 12 N 8B8B 7 Be 9C9C 10 C 10 B 11 N 11 B 9B9B 8 Be 20 Ne 22 Ne 21 Ne 9 Be 23 Na 17 Ne 16 F 15 F 22 Na 21 Na 20 Na 24 Al 23 Al 25 Al 24 Mg 23 Mg 22 Mg 21 Mg 20 Mg 19 Na (p,γ) (p,α) (β + ν) = studied at TAMU CNO, HCNO 25 Si 24 Si 26 Si June 2008 Comp with direct meas: 16 O( 3 He,d) 17 F vs. 16 O(p,  ) 17 F Gagliardi e.a. PRC 1999 vs. Morlock e.a. PRL Be(p,  ) 8 B (solar neutrinos probl.): p-transfer: S 17 (0)=18.2±1.7 eVb Breakup: S 17 (0)=18.7±1.9 eVb Direct meas: S 17 (0)=20.8±1.4 eVb Experiments using the ANCS

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

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

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

Rare Isotope Accelerators

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?

Radioactive Nuclei in Supernovae

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

Research and Education Well organized Institute People QualityInstruments Funding Support

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