1. Nuclear Reactions

2. Binding Energies The mass law below represents the masses of thousands of nuclei with a few parameters B=(Z(m p +m e )+(A-Z)m n - M(A,Z))c 2 Mass Excess  M= 9.31.478MeV (M(A,Z)-A) ; M in AMU Q value - energy released in exit channel of rxn assuming incoming kinetic energy small  M in -  M out B/A binding energy per nucleon

3. Nuclear Reactions Mass terms M(A,Z) = Zm p + (A-Z)m n m 1 = -a 1 A volume term m 2 = a 2 A 2/3 surface tension m 3 = a 3 (A/Z - Z) 2 /A symmetry term from Fermi energy of p&n Fermi-Dirac gases m 4 = a 4 Z 2 /A 1/3 Coulomb repulsion of protons m 5 =  (A) pairing energy - paired p or n more tightly bound set to find minimum in mass for a given A - valley of stability

4. Nuclear Reactions valley of stability - At high Z, nuclei are stable only if neutron # > proton # - coulomb term otherwise too large High Z elements neutron rich - initla stellar composition n poor - need rxns which are n sources

5. Nuclear Reactions The Coulomb barrier Classical limit –R nucleus ~ r 0 A 1/3 ; r 0 =1.2x10 -13 cm –r >> = h/mc x c/v QM limit – compton = h/mc = 1.13x10 -13 cm –for v/c ~ 0.25, ~ 4.5x10 -13 cm Rxn rate for flux of particles N p v into a target of area a, thickness x, and density N t

6. Nuclear Reactions In center of mass frame

7. Nuclear Reactions Assuming Boltzmann dist. Integrand max when  =E/kT+b/√E is a minimum gives shape of nuclear potential Coulomb part of potential nuc. pot.  v 2 /2

8. Nuclear Reactions Resonances After capture the new particle may be in an excited state of the compound nucleus. This increases the cross section for capture in a narrow energy range around the excited state with width  E= /  state Network equations A term exists for every possible rxn channel which creates or destroys j finite difference approx

9. Nuclear Reactions Terms such as Y j (t+  t)Y k (t+  t) go to Y j (t)  k+ Y k (t)  j +Y j (t)Y k (t) linearize - discard higher order terms in  An eqn linear in unknowns  can be written for each species The eqn for each species j contains a term  k for each species k connected to j by a rxn Write as a matrix  A=B where  is a column of  ’s A is a JxK matrix for J species with K terms - generally J=K with most entries = 0 B is a column of RHS rxn coefficients Y a Y b  N A Want  ’s =  Y to determine change in Y Solve  =BA -1 This formulation automatically includes reverse rates for rxns since for every matrix element j,k there is an element k,j which describes reverse rxn

10. Nuclear Reactions Nuclear rxns in stars can progress down three paths 1.Complete burning - most familiar H  He, He  CO ash is a minimum energy state 2.Steady state - dY i /dt=0 from contributions of several channels - CNO in H burning reach steady state abundances for a given T,  3.Equilibrium - forward/reverse rates balance. Get broad distribution of abundances determined by chemical potentials - minimize thermodynamic free energy of system Limiting rates determine speed of reaction - often weak interactions e.g. in PP chain 1 H(p,  + ) 2 D  ~10 9 yr

11. The Asymptotic Giant Branch When He core exhausted He shell burning begins Like H shell burning He shell drives the star redward - moves star along the Asymptotic Giant Branch roughly parallel to but higher in luminosity than the RGB

12. The Asymptotic Giant Branch When He core exhausted He shell burning begins Like H shell burning He shell drives the star redward - moves star along the Asymptotic Giant Branch roughly parallel to but higher in luminosity than the RGB Second dredge-up brings H burning products to surface H shell quenched until He shell moves out far enough to heat shell to burning T # of stars on AGB/# of stars on HB gives constraint on amount of time star spends in core He burning

13. The Asymptotic Giant Branch Extreme density gradients outside degenerate corre and burning shells

14. The Asymptotic Giant Branch Center of star is degenerate and cooling from weak emission - peak T not in core

15. The Asymptotic Giant Branch Star has extremely compact core - most of radius is extended envelope

16. The Asymptotic Giant Branch Star has extremely compact core - most of radius is extended envelope

17. The Asymptotic Giant Branch Double shell burning or Thermal Pulse AGB q(He) ~ 0.1q(H) so He shell catches up to H shell As He shell approaches H shell material expands, H shell quenched He burns outward, runs out of fuel, quenched H shell restarts, eats outward, ash builds up He shell ignites, repeat

18. The Asymptotic Giant Branch Double shell burning or Thermal Pulse AGB During He shell phases envelope convection penetrates deeply into star He shell produces small convective shell Non-convective mixing allows transport between shells mixing 12 C into H flame zone or p into He flame gives 12 C(p,  ) 13 N(  + decay) 13 C 13 C( ,n) 16 O is a neutron source - only works when p and He burning can mix

19. The Asymptotic Giant Branch Double shell burning or Thermal Pulse AGB s-process - slow n capture onto Fe peak seed nuclei - each n captured has time to  decay to a proton, increasing Z s-process takes place in intershell region where n produced primarily in intermediate mass stars just above maximum mass for He flash Produces species with A>90 3 rd dredge-up (actually numerous dredge-ups for each thermal pulse cycle) brings partial He burning products to surface with s-process enhancements - most efficient at low metallicity - C stars, Sr stars

20. The Asymptotic Giant Branch Double shell burning or Thermal Pulse AGB Produces species with A>90 s-process peaks where p & n form closed shells - p and n magic numbers i.e. 208 Pb with Z=82, N=126, both magic numbers even Z and even A nuclei more abundant

21. AGB Mass Loss

22. Often highly asymmetric (bipolar) AGB stars generally very cool - spectra dominated by molecular species –H 2 O, TiO, VO, Sr, Ba compounds, Si 2 O 3 ; SiC, C 2, Buckyballs in carbon stars Complex molecular spectra and low T allow line blanketing - much of high L goes into accelerating wind Atmospheres of cool stars dust rich –winds from direct radiation pressure –dust formation region can act like  mechanism - drive pulsations which become non-linear and create shocks in low  stellar atmosphere

23. AGB Mass Loss Thermal pulses during double shell burning can drive mass loss episodes Shell flashes - if H or He shell is degenerate when it ignites small explosion drives mass loss, may revivify proto-WD as red giant (Sakurai’s object) Small envelope above a burning shell can be removed in a short event - planetary nebula Fast wind from proto-WD evacuates bubble, causes Rayleigh- Taylor instabilities in swept-up shell add ionizing radiation from central star and get planetary nebula Low mass stars have only compact ionized bubble, high mass disperse envelope very quickly - only intermediate masses have visible PN with lifetime ~ 10,000yr

24. Morphology of Planetary Nebulae Many PN/proto-PN strongly bipolar IR and polarization show thick dusty torus Some axisymmetry due to rotation Mechanism for tight collimation unknown - B fields or companions possible Fliers - Fast, low ionization emission regions - clumps moving at hundreds of km s -1 near symmetry axis - mechanism unknown

25. Morphology of Planetary Nebulae Clumping - Shell of swept-up material breaks into dense clumps (n~10 4-6 cm -3 ) Two possible mechanisms - Rayleigh-Taylor instability from fast, low density wind impacting shell Or thermal instabilities - rapid efficient cooling ahead of shock causes fragmentation on scale where sound crossing time = cooling time