Evolution and nucleosynthesis prior to the AGB phase

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1 Evolution and nucleosynthesis prior to the AGB phase Amanda Karakas Research School of Astronomy & Astrophysics Mount Stromlo Observatory

2 Lecture Outline 1. Introduction to AGB stars, and the evolution prior to the AGB phase 2. Evolution and nucleosynthesis prior to the AGB phase 3. Evolution and nucleosynthesis of AGB stars 4. The slow-neutron capture process in AGB stars 5. Low and zero-metallicity AGB evolution 6. Super-AGB stars and post-agb objects

3 Outline of this lecture 1. Review of hydrogen and helium burning 2. Nucleosynthesis during core hydrogen burning 3. Surface abundance changes due to the first and second dredge-up events 4. The core He-flash 5. Nucleosynthesis during core helium burning

4 Central Burning lifetimes 25Msun 5Msun Burning process Temperature (Kelvin) Lifetime (years) Lifetime (years) H burning 3.8 x 10 7 (25) 6.7 x x x 10 7 (5) He burning 2.0 x 10 8 (25) 0.84 x x x 10 8 (5) Carbon burning 8.4 x Neon burning 1.6 x Oxygen burning 2.1 x Silicon burning 3.7 x days -- From Woosley, Heger & Weaver (2002, Rev. Mod. Phys. 74, 1015)

5 Energy content of burning processes Consider the reaction 12 C(α,γ) 16 O Entrance channel contains 12 C and 4 He; the exit channel contains 16 O and a gamma-ray photon Q-value is the energy released in the exit channel Define the atomic mass excess M az = (M az - A), in MeV is the rest-mass of 1 atomic mass units (amu) in MeV, M az mass of species (A,Z) Define Q, Q = c 2 ( M a + M X - M b - M Y ) = M a + M X - M b - M Y = ( ) = MeV Note that Q for H-burning ~ Mev for 4p to 4 He

6 Energy release per burning stage Define q = Q N a / A, where A = the total mass of nucleons q is the energy released upon consumption of a unit mass of fuel by the process in question 1 MeV/nucleon = X erg/g Process 4H 4 He 3α 12 C 4α 16 O 2 12 C 24 Mg 2 20 Ne 16 O + 24 Mg 2 16 O 32 S 28 Si 56 Fe q ( erg/g) 5 to to 0.3 q (MeV/nucleon) 5 to to 0.31

7 Population I clump giants The number of stars is proportional to the lifetime! Core He-burning clump Core H-burning From Faulkner & Cannon, 1973, ApJ, 180, 435

8 Hydrogen burning reactions Proton-proton chains Most important energy generating reactions in stars with M < 1.5 M sun (our Sun!) In the Sun, 86% of the time is spent in the PPI chain, where p + p D + e + + ν MeV D + p 3 He + γ (5.49 MeV) 3 He + 3 He 4 He + 2 p MeV The remainder of the time are spent in PPII (14%) and PPIII (0.02%) Energetic neutrinos from PPIII and other reactions have been detected, leading to the famous solar neutrino problem Where only 1/4 the number of ν s expected were detected Solution lies in neutrino physics. For more information:

9 Proton proton chains From MPA, Neutrino Astrophysics Group

10 PP chains Energy generation is proportional to ~T 4, compared to ~T 17 for CNO cycles

11 CNO cycles C, N, O act as catalysts for H-burning The number of C+N+O nuclei is conserved I.e. remains constant Nett result is 4 H 4 He C/N/O ratios in Sun (at the solar surface) and at the various CNO cycle equilibriums are very different! Effectively, all C+N+O is converted to 14 N The CN cycle operates first and at lower temperature to the ON cycles Nearly all of the energy comes from the CN cycle Temperature dependence of CN cycle is ~T 17

12 CN CNO cycles

13 CNO equilibrium ratios Ratios 12 C/ 14 N/ 16 O 12 C/ 13 C Surface of Sun 3/1/9 90 CNO equilibrium 1/120/10 ~3 C/N/O ratios at stellar surface and from the CNO cycle equilibriums are very different! 13 C and 14 N are enhanced 16 O abundance barely changed Low C-isotopic ratios at the surface of a star an indication that material was exposed to CN cycling

14 Advanced H-burning cycles

15 Helium burning Typical T ~ 1.5 x 10 8 K, density ~ 10 3 g cm -3 4 He + 4 He 8 Be MeV 8 Be is very unstable, decays in ~10-15 seconds But 8 Be + 4 He 12 C ** (7.65 MeV)(γ γ) 12 C (ground state) has a very large cross section due to a resonance near the Gamow peak (Hoyle 1954) Pair of reactions is the triple-α process, effectively 3 α 12 C Energy generation from the triple- α process is a steep function of T and is proportional to ~T 40 At slightly higher T, density the 12 C(α, γ) 16 O occurs At the end of core He-burning, the composition is ~50% 12 C and ~50% 16 O (depends on rates!) Values different for shell burning because higher T

16 Non-energetic reactions He-burning occurs in the ashes of H-burning Composition is 98% 4 He, ~2% 14 N Nitrogen-14 can capture alpha particles to produce secondary nuclei, depending on T: 14 N(α, γ) 18 F(β + ν) 18 O(α, γ) 22 Ne 22 Ne + α 25,26 Mg (+n, or γ) when T > 300 million K These may have important consequences, especially since the 22 Ne(α,n) 25 Mg reaction releases free neutrons that can be captured by Fe-group elements The other main neutron producing reaction, 13 C(a,n) 16 O will be discussed in the context of the slow-neutron capture process (Lecture 3)

17 Evolution as a function of mass Single stars with initial masses below ~8 solar masses do not proceed beyond core helium burning Neutrino emission processes cool the centre and prevent T becoming large enough to ignite carbon to burn (~800 million K) Instead H and He burn in shells around the CO core (will become the white dwarf) Stars in this phase of evolutionary phase are called asymptotic giant branch stars, or AGB stars Stars in the range 8 to 11Msun ignite carbon under degenerate conditions (similar to core He-flash) and are known as super-agb stars (Lecture 6)

18 Where mixing takes place TDU, HBB SDU FDU

19 First and second dredge-up Mixing episodes occur when the star becomes a red giant following core H or He-burning The stellar envelope becomes convective and eats down into the star as it becomes cooler, dredging up material processed during core H burning Changes to the surface composition a?er the first and second dredge-up (FDU, SDU) involve H- burning products Increases in 3,4 He, 13 C, 14 N, 23 Na and decreases in H, 12 C, 15 N, 16 O While convection moves inward for all stars a?er core He-burning, only M > 3Msun show changes to their surface composition and experience the SDU

20 Depth of the first and second DUPs Dashed-lines: SDU Solid lines: FDU

21 The second dredge-up: 6.5Msun Convective encompasses ~1.82Msun (0.28 x 6.5) during core H-burning! During FDU, envelope reaches down to ~1.35Msun, whereas during the SDU, the envelope reaches 0.95Msun

22 The first dredge-up Next five diagrams from my thesis. Composition profile as a function of mass, just a?er core H burning has finished:

23 FDU as a function of mass In the 1Msun, about 75% of the star is mixed by FDU In the 3Msun, about 85% of the star is mixed by FDU

24 FDU as a function of metallicity, Z In the 1Msun, about 75% of the star is mixed by FDU In the low-z 1Msun, about 72% of the star is mixed by FDU

25 Carbon isotope ratio A?er the FDU A?er the SDU

26 Nitrogen isotope ratio

27 O-isotopic ratios

28 Sodium

29 The need for extra-mixing Observations of C,N,O elements in evolved giants have allowed us to verify the accuracy of the first dredge-up models (e.g. Charbonnel 1994) In a large fraction of low-mass giants, the 12 C/ 13 C and 12 C/ 14 N ratios appeared lower than the predictions Different mechanisms have been proposed to allow for deviations from standard evolution theory (recently thermohaline mixing, Charbonnel & Zahn; Eggleton et al.) Some extra-mixing below the deepest extent that the FDU reaches, is needed for m < 2 Msun

30 Luminosity bump When the convection retreats a?er FDU it leaves behind a discontinuity in the abundance profiles This inhibits further mixing The H-shell eventually catches up and erases the discontinuity:

31 Extra mixing a?er FDU from Charbonnel (1994) Fig 3 from Charbonnel & Zahn (2007)

32 The core helium flash He is ignited under electron-degenerate conditions, with the He-luminosity reaching ~10 9 Lsun Takes about 10 6 years to remove degeneracy, and to move to quiescent He-burning During the flash, about 3% of the He is converted to 12 C The huge luminosities produced by the flash drive a convective region between the centre and base of the envelope Some debate if any mixing between the flash (which makes carbon) and the envelope could take place Unlikely except in Z = 0 stars (Lecture 5) Stars over ~2.5Msun do not experience the flash, because their cores do not become electron degenerate

33 Nucleosynthesis during the flash Flash driven convective pocket Envelope, surface composition About 3% of the He is converted to 12 C

34 Nucleosynthesis during He-burning The nucleosynthesis products from central He-burning do not reach the surface Exceptions: when binary interactions occur e.g. white dwarfs that explode as Type Ia supernovae, classical novae explosions For single stars, the He-burnt layers can escape from massive WR stars when they lose mass, or through the supernovae explosion He-burning is important because it determines the mass of the He-exhausted core and the composition of the eventual white dwarf The latter depends on the uncertain 12 C(a,γ) 16 O reaction, which determines the final C/O ratio in the core The mass of the core depends on semiconvective mixing and the inclusion of overshoot - both increase the mass

35 Core He-burning: 5Msun As 4 He burns it initially produces 12 C. But once there is a substantial amount of 12 C then we get 16 O from 12 C(α,γ) 16 O as shown in the central abundances of 12 C and 16 O, below:

36 End of core He-burning CO core He-burning shell 12 C 4 He H 16 O Results for a 6.5Msun, Z = 0.02 model: Central H-burning lifetime ~ 43 Myr Central He-burning lifetime ~ 9.6 Myr envelope

37 Convective overshoot We define a region to be convective if it satisfied the Schwarzschild criteria for convection That is, if in a region of the star, the ratio of the adiabatic to radiative temperature gradient is > 1, blobs of material are unstable to convective motions At this formal border the acceleration of material is zero but the velocity can be finite This is one reason why it is thought that overshoot, that is mixing beyond the formal boundary, takes place Rotation may also cause mixing beyond the formal boundary If overshoot is included then core H, He-burning lifetimes are increased because more fuel is being added The mass of the H, He-exhausted cores also increase

38 Effect of overshoot on HR diagram From Castellani et al. (2003)

39 Semiconvection During core H, He burning in massive stars (8Msun and above), and during He-burning in lower mass stars, we find semiconvection Where the concentration of chemical species are redistributed within a region that is stable to the standard Schwarzschild criteria for convection, but unstable to the Ledoux criteria A thermal gradient favours convection (Schwarzschild criteria) and relatively quick mixing time-scales but the competing effect of a molecular weight gradient (Ledoux criteria) will inhibit the mixing Main effect of semiconvection is to extend the core burning lifetime, and to increase the mass of the core

40 22Msun star from Heger et al. (2004)

41 Summary of 2 nd lecture Low and intermediate-mass stars go through central H and He-burning before reaching the AGB Mixing episodes occur, where the convective envelope moves inward, in mass, to where partial or full hydrogen burning took place during the main sequence The first dredge-up results in important surface compositional changes for low-mass stars with M < 3Msun Extra-mixing episodes may also occur in low-mass stars The SDU is the most important for more massive stars, between 4 to 8Msun Nitrogen increases at the surface, along with decreases in the carbon isotopic ratio He is ignited under degenerate conditions in stars less than about 2.5Msun, and non-degenerately in more massive objects

42 Summary of 2 nd lecture Overshoot and semiconvection increase the lifetimes and masses of core H, He-burning regions However the products of central He-burning, including the explosive core He-flash, do not reach the stellar surface Except in in a few cases that won t be discussed here including Type Ia supernovae and novae Massive stars over 8Msun also may release the products of He-burning to the interstellar medium when they become WR stars The core He-flash in Z=0 stars may drive mixing between the convective region and surface, see Lecture 5!

Evolution and nucleosynthesis of AGB stars

Evolution and nucleosynthesis of AGB stars Amanda Karakas Research School of Astronomy & Astrophysics Mount Stromlo Observatory Lecture Outline 1. Introduction to AGB stars; evolution prior to the AGB