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Dr. Amanda Karakas and Prof. John Lattanzio School of Physics & Astronomy Monash University Outline of Lectures The two lectures will be broken down into the following components: 1. Introduction – some basics 2. Evolution and nucleosynthesis of low and intermediate-mass stars 3. Super-AGB stars 4. The slow neutron capture process 5. Stellar yields Outline of Lectures The two lectures will be broken down into the following components: 1. Introduction – some basics 2. Evolution and nucleosynthesis of low and intermediate-mass stars 3. Super-AGB stars 4. The slow neutron capture process 5. Stellar yields Introduction • We are all made of star stuff! • 13.7 billion years ago, the big bang made (mostly) hydrogen (~76%) and helium (24%) • Everything else, including the material that makes up you and me, was cooked inside a giant stellar furnace • What stars produce what elements? • Our Sun is a star! How will it age? What elements will it make? Credit: HST public archive The Origin of the Elements Abundance distribution in our solar system 12 Hydrogen is 12 by definition Elements up to the Fe peak Neutron-capture elements 10 log10 (X/H) + 12 Carbon (Z=6) Iron (Z=26) 8 6 The aim of nucleosynthesis is to explain all of this graph! Rubidium (Z=37) 4 Barium (Z=56) 2 Lead (Z=82) 0 0 10 20 30 40 50 Atomic Number, Z 60 70 80 Using solar abundances from Asplund et al. (2009) Some basics • Low-mass stars: – Initial masses from 0.8 to ~2.0 solar masses • Intermediate-mass stars: – Initial masses from ~2.0 to ~8 – 10 solar masses • Massive stars: – Initial masses from 10Msun and up • These definitions for Z = 0.014 (solar); depend on Z • X = hydrogen mass fraction, Y = helium mass fraction, and Z = 1 - X - Y = “metals” • In the Sun: X = 0.7154, Y = 0.2703, Z = 0.0142 (Asplund et al. 2009) • [X/Y] = log10 (X/Y)star - log10 (X/Y)sun ; in our Sun [Fe/H] = 0.0 by definition Periodic Table Chart of the Nuclides The further a nucleus is from the valley of nuclear stability, the more unstable it is to β± decay i.e. the shorter is its half-life Credit: Frank Timmes Nucleosynthesis • Production of atomic nuclei in the Universe • Nucleosynthesis* takes place deep inside stars • How does it get out? Into the interstellar medium We need a way of moving the material from the stellar core, where thermonuclear reactions take place, to the surface. From there, we then need a way of moving the processed material into the ISM: 1. Low and intermediate-mass stars: Mixing and mass loss 2. Massive stars: mixing, mass loss and core collapse supernova explosions 3. Binary evolution channels: mass loss, envelope ejection, Type Ia supernova and nova explosions Star birth masses From Frank Timmes website The stars that make elements Very low mass stars (≤ 0.8Msun, depending on Z) are still on the main sequence fusing hydrogen in their cores ! These stars have not contributed to the chemical evolution of our Galaxy In terms of single stars, the most important are 1. Massive stars that explode as Type II (core collapse) supernova (> 10 solar masses); 2. Stars that evolve through the first and asymptotic giant branches (< 10 solar masses) How long do stars live? Age of the galaxy ≈ 12 x 109 years; Universe ≈ 13.7 x 109 years Initial mass (Msun) ! Main sequence lifetime! Total stellar lifetime! 25! 6.7 Myr! 7.5 Myr! 15 ! 11Myr! 13 Myr! 5! 80 Myr! 100 Myr! 2! 900 Myr! 1.2 Gyr! 1! 10 Gyr! 12 Gyr! 0.8! 20 Gyr! > 32 Gyr! 1 Myr = 1,000,00 years; 1 Gyr = 1000 Myr! Ages from Karakas & Lattanzio (2007); Woolsley et al. (2002)! How do stars work? • Once the temperature of a proto-star reaches about 10 million degrees Kelvin, nuclear fusion begins! • Hydrogen fusion or burning (i.e., similar to a H-bomb) • 4 protons ➙ 4He + 2 e+ + 2 neutrinos + energy • Where does the energy come from? • From E = mc2: the mass of 4 protons > 1 4He nuclei • How much energy? • Energy released = 25 MeV = 4 x 10-12 Joules But the Sun does it 1038 times a second! Into the interstellar medium We need a way of moving the material from the stellar core, where thermonuclear reactions take place, to the surface. From there, we then need a way of moving the processed material into the ISM: 1. Low and intermediate-mass stars (Karakas, Lattanzio) ! no explosion! Mixing + mass loss returns the material 2. Massive stars and core collapse supernova explosions (Chieffi, Müller) 3. Binary evolution, chemical enrichment and Galactic archaeology (Eldridge, Siess, Yong) Outline of Lectures The two lectures will be broken down into the following components: 1. Introduction – some basics 2. Evolution and nucleosynthesis of low and intermediate-mass stars 3. Super-AGB stars 4. The slow neutron capture process 5. Stellar yields Where mixing takes place Third dredge-up Hot bottom burning Second dredge-up First dredge-up Basic Stellar Evolution The x-axis: logarithm of the surface temperature of the star The y-axis: logarithm of the luminosity or amount of energy a star radiates per unit time. Main sequence: H to Helium τ ~ 1010 yrs for 1 ~ 108 yrs for 5 Red Giant Branch core contracts outer layers expand Second-giant phase after core He-burning star becomes a red giant for the second time Core He burning: He to Carbon Core H-burning: 1Msun Movies from John Lattanzio’s website: http://www.maths.monash.edu.au/~johnl/StellarEvolnV1/ Core H-burning: 5Msun Evolution after core H-burning • After core H-burning has ceased, the envelope expands and the core begins to contract • A hydrogen-shell burning is established in a shell around the contracting He-core • This provides most of the surface luminosity • At this point (owing to L = 4πσR2Teff 4), Teff drops because of increases to L and R • The envelope becomes convective, and moves inward into regions partially processed by previous H-burning ! This is the first important mixing event ! Known as the “first dredge-up” Up to the tip of the first giant branch • Envelope convection deepens • Mixes up material partially processed during the previous main sequence • This is the first dredge-up (FDU) • Main changes: – Reduction in Li, 12C/13C ratio – Increases in 3He, N – Oxygen isotopic ratios altered, depending on initial mass • Core He-burning ignited at tip of the giant branch • Mass loss will erode up to 20% of the envelope in ~1Msun stars Example: 1.25Msun, Z = 0.02 model Mass at tip: 1.12Msun Luminosity bump FDU ! Mass loss on RGB may not be this efficient, according to latest Kepler data of old open clusters (Miglio et al. 2012) The first dredge-up: 1Msun First and second dredge-up Dashed-lines: SDU Solid lines: FDU H-burning: Proton proton chains From MPA, Neutrino Astrophysics Group H-burning: CNO cycles CN H-burning: CNO equilibrium ratios Ratios Surface of Sun CNO equilibrium 12C/14N/16O 3/1/9 1/120/10 12C/13C 90 ~3 • C/N/O ratios at stellar surface and from the CNO cycle equilibriums are very different! • 13C and 14N increase • 16O abundance barely changed • Low 12C/13C ratios (< 20) at the surface of a star an indication that material was exposed to CN cycling Advanced H-burning cycles • Increase in the abundance of 23Na • Produces radioactive 26Al and re-arranges the heavy Mg isotopes What do we see at the surface? 30 after FDU after SDU 12 C/13C ratio 28 26 24 22 Initial value = 87 20 18 1 2 3 4 5 6 Initial stellar mass (Msun) From Karakas & Lattanzio (2014) 7 8 Oxygen isotope ratios after FDU • Mixes up material partially processed during the previous main sequence • Oxygen isotope ratios: 16O/17O can decrease by up to 80% whereas 16O/ 18O increases by ~30% (e.g., Boothroyd & Sackmann 1997) 3000 660 Initial value = 2630 Z=Zsolar O/18O ratio 2000 640 1500 16 O/17O ratio 2500 16 680 after FDU after SDU 1000 620 600 580 Z=Zsolar 560 500 540 520 0 1 2 3 4 5 Initial stellar mass (Msun) 6 16O/17O 7 ratio at surface after FDU (red) and SDU (blue) From Karakas & Lattanzio (2007) Initial value = 500 1 2 16O/18O after FDU after SDU 3 4 5 Initial stellar mass (Msun) ratio at surface after FDU (red) and SDU (blue) 6 7 First and second dredge-up: Issues Observations of low-mass giants in the field or in clusters show that the 12C/13C ratio is less than 20 (e.g., Gilroy 1989) 30 after FDU after SDU 12 Predicted range C/13C ratio 28 26 24 22 Initial value = 87 20 Observational range 18 1 2 3 4 5 6 Initial stellar mass (Msun) 7 8 Extra mixing in low-mass giant stars • • • • • • • • M < 2Msun Standard stellar models: Only one mixing event between MS and tip of the first giant branch The first dredge-up: 12C/13C ~ 20, C/N ~ 1.5 Disk FGB stars (e.g., Gilroy 1989) have 12C/13C ~ 10, and C/ N ~ 1.0 Evidence that some form of chemical transport is acting in low-mass giant envelopes Mechanism? Thermohaline mixing currently favoured Charbonnel & Zahn (2007) See also Eggleton et al. (2008), Stancliffe et al. (2009), Angelou et al. (2012), Lattanzio et al. (2015) Core helium burning • The helium core continues to contract and heat • Once the temperature inside the core reaches about 108 K, core He ignition takes place • Low-mass stars (< 2.0Msun) need to contract substantially before reaching this temperature ! the central regions to become electron-degenerate • Neutrino energy losses from the core cause the temperature maximum to move outward • Eventually, the triple alpha reactions are ignited at the point of maximum temperature • E.O.S only slightly dependent on T, leading to a thermonuclear runaway: The core He flash The core helium-flash The core helium-flash Core helium burning • Following core He-ignition, there is a stable period of core helium fusion • The coulomb repulsion is larger for He than for H • More energy is required to fusion to occur • This means higher burning temperatures and because energy generation ∝ T40, shorter lifetimes! • Typical He-burning lifetimes are ~100 million years for low-mass stars (~1Msun), compared to 1010 for H-burning • Whereas core He-burning lasts about 20 million years for the 5Msun, compared to 80 million years for H-burning Helium burning • Simultaneous collision of 3 4He nuclei too rare to provide the burning rate required • Fred Hoyle (1954) predicted a resonance in 12C to speed up the collision • Hoyle’s state was experimentally measured shortly after his prediction • Typical T ~ 1-2 x 108 K, density ~ 103-4 g cm-3 • The main reaction takes place in three steps: Step 1: 4 Step 2 : 4 Step3 : 12 4 8 8 12 He+ He ↔ Be He+ Be ↔ C * C * → 12C + 2γ • The 12C* indicates that the nucleus is in an excited state Helium burning • At slightly higher T and density the 12C(α, γ)16O reaction occurs once a supply of 12C is available • The 12C(α, γ)16O reaction also supplies energy • At the end of core He-burning, the composition of the core is roughly 50% 12C and 50% 16O • Although the final C/O greatly depending on the rates and can be as extreme as C:O = 0.10:0.90. • Temperature dependence for the triple-α rate turns out to be roughly ε ∝ T40! • This means that helium burning leads to a very steep temperature gradient ! convection Basic Stellar Evolution The x-axis: logarithm of the surface temperature of the star The y-axis: logarithm of the luminosity or amount of energy a star radiates per unit time. Main sequence: H to Helium τ ~ 1010 yrs for 1 ~ 108 yrs for 5 Red Giant Branch core contracts outer layers expand Second-giant phase after core He-burning star becomes a red giant for the second time Core He burning: He to Carbon After core helium burning • Following the exhaustion of helium, the core contracts and heats • The composition is roughly 50% carbon and 50% oxygen, • The core composition depends on the rates of the triplealpha (3 helium nuclei ! 12C atom) and 12C(α,γ)16O reactions • The later reaction is quite uncertain and difficult to determine in a laboratory • The mass of the core depends on convection, which is a big uncertainty in stellar evolution models The early asymptotic giant branch • Following core He-exhaustion, the star evolves up the second giant branch, or AGB • A helium burning shell is established around the contracting C-O core, which narrows as the star evolves • Eventually the shell becomes thin and partially degenerate • Helium burning is unstable under such conditions ! leads to thermal pulses or He-shell instabilities • However, the early part of the AGB is the longest in time and is where the second mixing event occurs Structure during second dredge-up Results for a 5 Msun, Z = 0.02 model: The second dredge-up: 5Msun Where mixing takes place Third dredge-up Hot bottom burning Second dredge-up First dredge-up Mass loss • Convective mixing (dredge-up) can mix the products of nucleosynthesis from the (hot) interior to the surface. • Mass loss removes the enriched envelope, expelling the products into the interstellar medium. ! When does most of the mass loss occur? For low and intermediate-mass stars, that is during the asymptotic giant branch (AGB) ** caveat: except in the case of the 1Msun Asymptotic Giant Branch stars • The asymptotic giant branch is the last nuclear burning phase for stars with mass < 8Msun • AGB stars are cool (~3000 K) evolved giants, spectral types M, S, C H-rich envelope • It is during the AGB where the products of nucleosynthesis reach the stellar surface • Many AGB stars are observed to be losing mass in dense outflows of material " Enriching the interstellar medium " Progenitors of planetary nebulae " Review by Herwig (2005, ARAA) and Karakas & Lattanzio (2014, PASA) H-exhausted core Products of nucleosynthesis Low and intermediate-mass stars go through central hydrogen and helium burning During the AGB, they have shells burning H and He 1. First dredge-up: Products of (partial) H burning 2. Second dredge-up: Products of H burning 3. Third dredge-up: Products of H, He-burning and neutroncapture nucleosynthesis 4. Hot bottom burning: Products of H-burning 5. Extra mixing processes: Products of H-burning* ! We we will now discuss the AGB phase of evolution . AGB nucleosynthesis 4He, 12C, 19F, s-process elements: Zr, Ba, ... At the stellar surface: C>O, products of Heburning Envelope burning: 4He, 14N, 23Na Interpulse phase (t ~ 102-5 years) He-shell instabilities • The He-shell thins as the star ascends the AGB and becomes thermally unstable • He-burning in a thin shell leads to a thermal runaway, similar to the core He-flash • Why? • Not caused by electron degeneracy, although the shell is partially degenerate • Caused by the shell being thin! • Contracting shell → hotter → ε∝ T40 → but shell can’t expand enough to cool → thermal runaway • Luminosities can reach > 108 solar luminosities He-shell burning in AGB stars • Up to ~108 Lsun can be generated by a thermal pulse • Energy goes into expanding the star • He-shell becomes unstable to convection ! mixes products of Heburning throughout shell 6Msun, Z = 0.014 model star: Intershell convection during TPs • The enormous amount of energy drives a convective region in the intershell • Because energy too much for radiation to carry • Extends over almost the whole intershell • Homogenises abundances within this region • The mass of the pocket ~ few 10-2 Msun, depending on the stellar mass • The duration of convection is ~few hundred years Convection zones = green, radiative = pink Results for a 1.9Msun, Z = 0.008 model Model number proxy for time The thermal pulse cycle Deep convective envelope H-burning shell He-rich intershell He-burning shell Carbon-Oxygen core The AGB Evolution Cycle 1. 2. 3. 4. On phase: He-shell burns brightly, producing up to 108 Lsun, drives a convection zone in the He-rich intershell and lasts for ~ 100 years Power-down: He-shell dies down, energy released by flash drives expansion which extinguishes the H-shell Third dredge-up: convective envelope moves inward into regions mixed by flash-driven convection. Mixes partially He-burnt material to surface. Interpulse: star contracts and H-shell is re-ignited, provides most of the surface luminosity for the next ~105 years Pulse (He-burning) ! TDU (mixing) ! Interpulse Few ~102 yrs ! ~102 years ! ~105 yrs Let’s look at a TP again Extent of convective pocket is 1.7 x 10-2 Msun About half gets mixed into envelope Mass of H-exhausted core is decreased by convection moving inwards He-exhausted core 22nd thermal pulse for the 3Msun, Z = 0.02 model What is all the fuss with the TDU? • The third dredge-up determines how much He-shell material is mixed from the core to envelope • That is, it determines the role that AGB stars play in the evolution and origin of elements in the Universe!! Third dredge-up • Badly named, can re-occur after each thermal pulse • Inward movement of convective envelope, reaches into the He-shell • Right-hand panel shows the evolution of the core in a low-mass AGB model • Six (third)-dredge-up events are visible. Each one will mix He-shell material to the surface H-rich envelope He-rich intershell C-O core Typical Galactic C-rich AGB star: 1.8Msun, Z = 0.01 Non-energetic reactions • He-burning occurs in the ashes of H-burning • The composition is typically 98% 4He, ~2% 14N • Remember that the CNO cycle produces mostly 14N, which can capture alpha particles to produce secondary nuclei, depending on T: – – 14N(α, γ)18F(β+ν)18O(α, γ)22Ne 22Ne + α → 25,26Mg (+n or γ) when T > 300 million K • These reactions produce little energy but are important for nucleosynthesis • Example, the 22Ne(α,n)25Mg (Q = -0.478MeV) reaction releases free neutrons that can be used to produce heavy elements i.e., 56Fe(n,γ)57Fe(n ,γ)… Fluorine production • It’s complicated! • The reaction chain: 18O(p, α)15N(α, γ)19F(α, p)22Ne • Fluorine production takes place in the He-intershell: This is a region rich in 4He, 12C • There are almost no protons or 15N • These are created by other reactions including: – – – – – – – γ)18F(β+)18O - main reaction to produce 18O 13C(α, n)16O - produces free neutrons (also for the s-process) 14N(n, p)14C - produces free protons 18F(α, p)21Ne - new, alternative proton production 14C(α, γ)18O - alternative reaction 18O(α, γ)22Ne - main 18O destruction reaction 15N(p, α)12C - destroys 15N 14N(α, Helium burning: summary From a nucleosynthesis point of view: • The triple alpha and 12C(α, γ)16O reactions convert 4He into 12C and 16O • Secondary reactions can produce 18O, 19F, 22Ne, 25Mg, 26Mg • Final composition depends on temperatures, densities, and the duration of burning • Secondary reactions can produce free neutrons (e.g., 13C(a,n)16O, 22Ne(a,n)25Mg) Products of He-shell nucleosynthesis 3Msun, Z = 0.014: Surface abundance of carbon (left) and fluorine (right) during the AGB ! We can make a carbon-rich star, which has C/O > 1 3.5 -5.6 C/O ratio at surface -5.7 3 log10 (19F) abundance -5.8 C/O ratio 2.5 2 C/O = 1.0 1.5 1 -5.9 -6 -6.1 -6.2 -6.3 -6.4 -6.5 0.5 -6.6 0 0 5 10 15 thermal pulse number 20 25 -6.7 0 5 10 15 thermal pulse number 20 25 Third dredge-up uncertainties • It is important to know if the models are providing an accurate description of mixing in real AGB stars • Because the third dredge-up determines how much Heshell material is mixed from the core to envelope • Do current models predict enough TDU? • Or too much? ! Do the model predict the right mass and luminosity ranges for carbon stars? Carbon star luminosity functions • Distances to the Magellanic Clouds are known • Can derive accurate C-star luminosity functions • These indicate that (most) stellar models do not predict enough dredge-up at low enough masses • And it is deeper at these lowest masses than current models predict • Can “force” the TDU in lowmass models… Binaries Stancliffe, Izzard, & Tout (2005) Uncertainties: The amount of third dredge up 0.57Msun 0.55Msun 1.25Msun, Z = 0.01: Forcing dredge-up by extending the base of the envelope by N scaleheights e.g., Karakas et al. (2010); Frost & Lattanzio (1996) 2Msun, [Fe/H]= -5.45: Diffusive mixing + Herwig’s scheme for extending the envelope using exponentially decaying overshoot From Simon Campbell Intermediate-mass AGB stars Along with thermal pulses and the third dredge-up, these stars also have: • Second dredge-up: Biggest ΔY (up to 0.1) • Hot bottom burning: Proton-capture nucleosynthesis at base of envelope (products: N, Na, Al) Log Y (=X/A) Example: 6Msun, Z = 0.02 Hot bottom burning & TDU Example: 6Msun, Z = 0.02 Third dredge-up (TDU) and HBB act together CN cycle is acting close to equilibrium for ~20 thermal pulses 20 (b) 0.5 14 C/O ratio C/13C ratio 16 12 6Msun, Z =0.02 0.6 Initial 12C/13C ratio = 86.5 18 12 10 8 0.4 0.3 0.2 6 0.1 4 2 0 5 10 12C/13C 15 20 25 30 Thermal pulse number 35 40 45 ~ 3 is the equilibrium ratio 0 0 5 10 15 20 25 30 Thermal pulse number 35 40 The C/O ratio never exceeds 1 45 Hot bottom burning & TDU Looking at the surface abundances of Ne to Al as a function of metallicity: • 6Msun, Z = 0.02 has a peak temperature of ~80 million K • 6Msun, Z = 0.004 has a peak of ~95 million K Lithium production • The first thing to happen is that 7Li is produced via the Cameron-Fowler Beryllium Transport Mechanism • This is basically PP chains plus convection! • The idea is that lithium is made by 3He(α, γ)7Be • and then to use convection to move the 7Be away from the hot region before it can complete the PPII or PPIII chains: BAD! BAD! Lithium production Lithium is produced by the Cameron-Fowler mechanism: 7Be is transported by convection, where it captures an electron to produce 7Li Log ε(Li)max = log10(Li/H)+12 = 4.5 6Msun, Z = 0.02 Uncertainties caused by convection Surface luminosity as a function of time for three convective prescriptions Surface CNO abundances as a function of total mass From Ventura & D’Antona (2005) Other mixing phenomena • What is the impact of non-convective extra mixing processes on AGB evolution and nucleosynthesis? • Examples include: rotation, thermohaline or double diffusive mixing, mixing induced by internal gravity waves, magnetic fields… • Thermohaline mixing: connects H-shell to base of envelope allowing the products of (partial) H-burning to reach the surface (e.g., 3He down, 13C up) I won’t have time to discuss all of this here! Reading: Karakas & Lattanzio (2014, PASA review, arXiv: 1405.0062) Outline of Lectures The two lectures will be broken down into the following components: 1. Introduction – some basics 2. Evolution and nucleosynthesis of low and intermediate-mass stars 3. Super-AGB stars 4. The slow neutron capture process 5. Stellar yields Carbon burning • Stars from 8Msun and above may experience core carbon burning • Many outcomes of 12C +12C are possible (20Ne, 23Na) • There are also many light-particle reactions on 12C, 16O, 23Na that occur at the same time • For hydrostatic C-burning: T ≈ 0.6 - 1.0 GK • For explosive C-burning: T ≈ 1.8 - 2.5 GK • Explosive burning occurs during the SN explosion, when the shock superheats the outlying material to very high temperature for a short time • Timescale for central C exhaustion: ~ few x 103 years, depending on T and density Off-centre carbon ignition • M > 8Msun (at Z = 0.02) up to 11Msun • These stars go through degenerate carbon ignition • Before ascending the thermallypulsing AGB with O-Ne cores • Q: What fraction explode as supernovae or leave massive white dwarfs? • E.g., Nomoto (1984), Poelarends et al. (2008) • The brightest AGB stars in young populations, with Mbol ~ −7.6, brighter than the traditional AGB limit (Mbol ~ −7.1) 7.5Msun, Z= 10-4 model by Siess (2007) Carbon ignition: 9Msun, Z = 0.02 Maximum temperature peaks at ~950 x 106 K. Duration of carbon flashes and central burning ~30,000 years (model from Karakas et al. 2012) Super-AGB stars A 9Msun, Z = 0.02 model has a core mass of ~1.18Msun. Too low to become an electron capture supernovae (from Karakas et al. 2012) It will produce an O-Ne white dwarf Nucleosynthesis in super-AGB stars Log Y (=X/A) 7Msun, Z = 0.002 (1/100th solar). Peak temperature ~ 140 x 106 K. This is about as extreme as it gets in an AGB star! Recent models by: Siess (2007, 2010), Pumo et al. (2008), Doherty et al. (2010), Karakas et al. (2012), Herwig et al. (2012), Takahashi et al. (2013) Doherty et al. (2014a,b) Doherty et al. (2015) Super-AGB stars The final fate of super-AGB stars is uncertain ! Will they mostly produce massive ONe white dwarfs ! What fraction will explode as electron capture supernova? ! What are their nucleosynthesis products? H burning? He-shell burning? The rapid neutron capture process? ! What happens when they are in a binary system? Will more explode? ! How do they affect the enrichment of the galaxy? Lots of questions! Very exciting stuff From Doherty et al. (2015) Outline of Lectures The two lectures will be broken down into the following components: 1. Introduction – some basics 2. Evolution and nucleosynthesis of low and intermediate-mass stars 3. Super-AGB stars 4. The slow neutron capture process 5. Stellar yields AGB stars as element factories Credit: IAC/ESA Production of heavy elements • By heavy elements we mean heavier than iron (Fe) • Z is large ! the electrostatic repulsion inhibits fusion reactions • Most heavy nuclei are formed by neutron addition onto Fepeak elements • Two processes: – r-process (rapid neutron capture) – s-process (slow neutron capture) Reading: Meyer (1994), Sneden, Cowan & Gallino (2008) The distribution of heavy elements s-process peaks Making heavy elements Rare proton-rich isotopes Proton number The radioactive Tc is observed in stars! Branching point Neutron number Magic Numbers • Analogous to the behaviour of atomic physics, where certain electron configurations are very stable • Such elements are inert (e.g., noble gases, He, Ne, Ar) • Similar behaviour seen in nuclear physics for isotopes with full neutron shells: neutron magic numbers • nuclei with a magic number of neutrons n = 50, 82, 126 (for lighter elements n = 2, 8, 20, 28) • A nuclei composed of a magic number of protons AND neutrons is very stable, doubly magic e.g. 16O, and 208Pb with p = 82 and n = 126 • Closed-shell nuclei are very stable against neutron capture and have low neutron-capture cross sections • Act as bottlenecks, seen as s-process peaks • The r-process peaks produced when very unstable nuclei decay to nuclei with a magic number of neutrons • Note that 56Ni, made in abundance by SN, is doubly magic! The slow neutron capture-process • Add neutrons slowly, so that unstable nuclei generally have time to β-decay before capturing another neutron • That is, τbeta-decay << τneutron-capture • Produces nuclei along the “valley of beta-stability” • Start with Fe-peak elements as they are very abundant, a typical path is: 56Fe(n,γ)57Fe(n,γ)58Fe(n,γ)59Fe • What happens at 59Fe? (half life ~ 44 days) ! branching point: Choice between decaying to 59Co or capturing a neutron to form 60Fe • The isotopic component of all elements between Sr and Pb have an important contribution from the s-process What stars make s-process elements? 1. AGB stars – AGB stars are observed to be rich in carbon and heavy elements produced by the s-process, including the unstable element Tc – The discovery of Tc in the atmosphere of an AGB star in 1952 was the first confirmation that stars make heavy elements 2. Massive stars – The He and C-burning shells of massive stars also make heavy elements via the s-process . AGB nucleosynthesis 4He, 12C, 19F, s-process elements: Zr, Ba, ... Interpulse phase (t ~ 102-5 years) At the stellar surface: C>O, sprocess elements The s-process in AGB stars • AGB stars observed to be enriched in s-process elements are mostly of low mass, between ~1 to 3Msun • How does the s-process occur? • The trick is making free neutrons! • These come from the13C(α, n)16O reaction • He-shell is ashes of H-burning: 98% He, 2% 14N, some 13C • Models have shown that the amount of 13C is not enough to account for s-process enhancement of S and C-type stars • We still do not understand how stars produce enough 13C in a He-burning region to activate this reaction Reading: Busso, Gallino & Wasserberg (1999), Käppeler et al. (2011), Karakas & Lattanzio (2014) Neutron production Neutrons form in 13C pockets – we don’t know how these form! 13C(α,n)16O Neutron production Extra burst of neutrons from the 22Ne(α,n)25Mg reaction, which takes place during thermal pulses 13C(α,n)16O 22Ne(α,n)25Mg Theoretical models Typical neutron density profile in time: Low mass Neutron source 13C(a,n)16O 22Ne(a,n)25Mg Maximum neutron density 108 n/cm3 1013 n/cm3 Timescale 10,000 yr Neutron exposure 0.3 mbarn-1 Intermediate mass 10 yr 0.02 mbarn-1 (at solar metallicity) Neutron sources: 22Ne source • • • • The other neutron source is 14N(α,γ)18F (β+ν)18O(α,γ)22Ne(α,n)25Mg Plenty of 14N left over from CNO cycling to produce the 22Ne Occurs in stars when the temperature of the He-burning region exceeds about 300 x 106 K, under convective conditions Intermediate-mass AGB or super-AGB stars (~4 to maybe 10Msun) The s-process in solar-metallicity models Z = 0.014, [Fe/H] = 0 Similar to the bulk metallicity of the disk of our Galaxy In models a factor of 2 less in metallicity Z = 0.007, [Fe/H] = -0.3 Similar to the bulk metallicity of the Large Magellanic Cloud The s-process: The effect of metallicity Decrease in metallicity results in more s-process elements at the 2nd peak (Ba, La), then at the 3rd (Pb) 3.5 3 2.5Msun, [Fe/H] = -1.4 2.5Msun, [Fe/H] = 0 2.5Msun, [Fe/H] = -2.3 2.5 [X/Fe] 2 1.5 1 0.5 0 Sr=38 Ba=56 Pb=82 -0.5 30 40 50 60 Atomic Number 70 e.g., Busso et al. (2001) 80 The s-process: Why range of mass? • The s-process in a 6Msun, Z = 0.0001 AGB star produces copious Rb (Z=37) compared to Ba, Pb • This is because it occurs at high neutron densities: ~1013 n/cm3 • Figure: Yields for [Fe/H] = −2.3 from Lugaro et al. (2012) for M = 1 to 6Msun 3.5 2Msun, [Fe/H] = -2.3 6Msun, [Fe/H] = -2.3 3 [X/Fe] Sr=38 Rb=37 2.5 2 1.5 1 0.5 0 Ba=56 -0.5 30 40 50 60 Atomic Number Pb=82 70 80 Neutron density indicators At high neutron densities, two branching points open that allow rubidium to be produced. At Nn=5 x 108 n/cm3 ~80% of the flux goes through 85Kr, and the branching at 86Rb opens to make 87Rb 1. 85Rb has a high σ = 240 mb (30 keV) 2. 87Rb is magic, has a low σ = 15 mb (30 keV) ! Rb in AGB stars in an indicator of the neutron density! 86Kr, 87Rb, and 88Sr are all magic, with low neutron capture cross sections In low-mass stars: 88Sr produced In massive AGB: 87Rb Zr,Sr/Rb > 1 ! low-mass AGB Zr,Sr/Rb < 1 ! IMS AGB Rubidium in bright AGB stars [Rb/Fe] Max. production factor ~ 100! Increasing stellar mass From D. A. Garcia-Hernandez et al. (2006, Science) Problems with theoretical picture Post-AGB stars: Evolved from stars of low-mass, 1-1.5Msun at relatively low metallicity, [Fe/H] ~ -1 (e.g., De Smedt et al. 2014; van Aarle et al. 2013) Figure from Kenneth De Smedt LMC object,[Fe/H] = -1 What about carbon-enhanced metalpoor stars? • [Pb/La] from a selection of CEMP stars. From Van Eck et al. (2003) 2Msun, [Fe/H] = -2.3 ! Produces CEMPtype composition ! Low neutron density 6Msun, [Fe/H] = -2.3 ! Does not produce a CEMP ! High neutron density Updated model predictions from Lugaro, Karakas et al. (2012) & Karakas (2010) Neutron production is still poorly understood Neutrons form in 13C pockets – we don’t know how these form! 13C(α,n)16O Neutron production is still poorly understood Neutrons form in 13C pockets – we don’t know how these form! What is the effect of rotation? Or of the sudden mixing of protons? 13C(α,n)16O Open questions • Rotation completely inhibits the s-process (Herwig et al. 2003) or moderates its effects (Piersanti et al. 2013) • However, there are very few models! • The abundance distribution of low-metallicity post-AGB stars and the composition of CEMP s/r stars point toward another process that occurs in nature ! intermediate-neutron capture process, or “i-process” ! Low-mass (< 1.5Msun) and/or super-AGB stars are proposed sites… ! How does the i-process contribute toward the enrichment of the Galaxy, if it does at all? Exciting times! Outline of Lectures The two lectures will be broken down into the following components: 1. Introduction – some basics 2. Evolution and nucleosynthesis of low and intermediate-mass stars 3. Super-AGB stars 4. The slow neutron capture process 5. Stellar yields Definition of a stellar yield General definition: Amount of matter (ΔM) that is expelled from a star over the course of its life More precisely… integrated amount (in mass) of species k that is expelled into the interstellar medium (ISM) over the stellar lifetime, τ, ∫τ [ Xk (t) - Xk(0)] (dM/dt) dt, minus the amount of k that is initially present in the wind (Xk(0) ΔM) How do we compute yields? • Calculate the evolution of a model star with a given initial mass (M) and metallicity (Z) • Use favourite mass-loss law • Evolve the model from the zero-age main sequence through all intermediate evolutionary stages to the tip of the TP-AGB • Stop when all the envelope has been removed by mass loss • Integrate the wind lost over the lifetime for each species in the nuclear network Problems with this picture 1. TP-AGB models can be very difficult to run (especially true for low metallicity and super-AGB stars) 2. Not always possible to evolve a stellar evolution code through all evolutionary phases (e.g. the core He-flash or final envelope ejection) We often have to make a few simplifications in practice, such as removing the remainder of the envelope when the calculations die, which may ~0.5Msun for a 6Msun star (what uncertainties will be caused by this?) AGB chemical yields Example: [Fe/H] = 0 (solar) 12C Yield = amount of an isotope ejected into the ISM over the star’s lifetime 14N • • 19F 17O M = 1 to 6Msun Metallicities [Fe/H] = 0, −0.4, −0.7, −2.3 What is lacking? • No s-process • No super-AGB stars • No non-standard stellar physics (e.g., rotation) From Karakas (2010) Yields and surface abundances for hydrogen to sulphur AGB chemical yields Light elements only (no s-process): • Siess (2010) and Doherty et al. (2014a,b) – yields for super-AGB stars • Lagarde et al. (2012) –large range of stellar masses, metallicities exploring non-standard physics (extra mixing and rotation) but only two AGB models; • Charbonnel & Lagarde (2010) – Lithium only • Ventura et al. (2013) – extended mass range (~1.5 to 6-8Msun) including super-AGB yields. Range of metallicities but maximum Z = 0.008 (aimed at globular cluster chemical evolution) Yields of s-process elements • FRUITY database: Cristallo et al. (2011) and Straniero et al. (2014) yields for 1-3Msun for a large range of metallicities; an increasing number of intermediate-mass models up to 6Msun • Our group: Lugaro et al. (2012) and Fishlock et al. (2014) yields of 1 to 6-7Msun for [Fe/H] = -2.3, -1.2 • NuGrid/MESA: Pignatari, Herwig et al. (2013, ApJ, submitted) for Z = 0.01 and 0.02 for limited masses • At very low metallicities: Campbell et al. (2010) and Cruz et al. (2013) but no tabulated yields What is lacking? Yields for low metallicity for all masses. Super-AGB yields. Full grid for solar metallicity Yields of s-process elements Let’s look at the yields from models of [Fe/H] = -1.2 From Fishlock, Karakas et al. (2014) Weighting by an initial mass function Z = 0.014, [Fe/H] = 0 The contribution of AGB stars Low and intermediate-mass stars are important factories for the following elements: • Li, C, N, F (e.g., Romano et al. 2010) • Neutron-rich isotopes of C, O, Ne, Mg, Si • And about half of all heavy elements including Sr, Y, Ba, Pb (e.g., Bisterzo et al. 2014) Evolution of elements in the solar neighbourhood From Kobayashi, Karakas, & Umeda (2011) Using AGB yields from Karakas (2010) and Campbell & Lattanzio (2008) C N F 16O/18O ratio at surface after FDU (red) and SDU (blue) Chemical evolution Spectroscopic observations of Galactic disk and halo stars at various metallicities The curves represent the total s+r process contribution for the different elements [Ba/Fe] [Eu/Fe] Halo - dotted lines Thick disk - dashed lines Thin disk - solid lines For the thin disk the AGB s-process contribution also show as thick solid line From Travaglio et al. (2004) [Ba/Eu] Summary of lectures • We’ve covered the evolution and nucleosynthesis of low and intermediate-mass stars • Looked at the production of elements from H and Heburning during the pre-AGB and AGB phase • And the production of heavy elements via the slow neutron capture process • Finally, we introduced the definition of a stellar yield and looked at how AGB stars may affect the enrichment of the Universe over time • I’ve also tried to include some open issues and questions, which will drive the future of this field