AST 1410 – Stars (2018)
Table of Contents
1 Syllabus
 Lectures
 M12 and W111 (i.e., 2 hr, with coffee break), AB 113
 Lecturer
 Marten van Kerkwijk, MP 1203B, 4169467288, mhvk@astro (.utoronto.ca)
 Office hours
 Drop by my office, or by appointment
 Web page

http://www.astro.utoronto.ca/~mhvk/AST1410/
 notes
Stellar astrophyiscs – the success story of 20th century astronomy – requires a synthesis of most of basic physics (thermodynamics, quantum mechanics, and nuclear physics). It underlies nearly all of astronomy, from reionisation to galaxy evolution, from interstellar matter to planets, and from supernovae and planetary nebulae to white dwarfs, neutron stars, and black holes.
In this course, we will review these successes (roughly first four weeks) and then discuss current topics and remaining puzzles (in four twoweek series, detailed content depending on interest).
Underlying Physics
 Master equations
 equilibria; timescales; massradius and massluminosity relations; HertzbrungRussell diagram, common threads in stellar evolution, features in stellar evolution.
 Equation of state
 fermions and bosons; pressure and energy density; ideal gas; (completely and partially) degenerate gas; radiation pressure; Boltzman distribution; Saha equation.
 Heat loss
 radiative diffusion; conduction; opacity sources; Schwarzschild and Ledoux criteria; mixing length theory; convective flux; stellar context for convection; semiconvection.
 Energy production
 nuclear binding energy; Coulomb barrier; reaction channels (PP, CNO, He and beyond, D/Li burning, s/r/pprocesses); rates and neutrinos.
Stellar Evolution Themes
 Lowmass stars
 Hayashi track; Li burning; (former) solar neutrino problem; pressure ionization and thermal ionization; convection zone advance; rotational evolution; shell burning; coremass radius, luminosity relations; helium core flash; thermal flashes; RGB/AGB winds; production of intermediatemass elements; white dwarfs.
 Highmass stars
 CNO burning and core convection; Eddington luminosity and formation/mass loss of high mass stars; nucleosynthetic yield of high mass stars; rotational evolution; feedback to the galaxy; electroncapture, core collapse and pair instability SNe; Pop III stars; neutron stars, black holes.
 Binary evolution
 frequency of binarity; tidal synchronization and circularization; Roche lobe overflow; conservative and nonconservative mass transfer; commonenvelope evolution; mergers; blue and red stragglers.
 Nucleosynthesis
 production of the elements; explosive nucleosynthesis; r and s process; corecollapse supernovae; thermonuclear supernovae; various dredgeups; thermal pulses.
Course texts
The main book we will use is Stellar Structure and Evolution (KW; Kippenhahn & Weigert for first, plus Weiss for second edition; SpringerVerlag, 1990, 2012). Especially for those who did not take undergraduate astrophysics, I strongly recommend An Introduction to Modern Astrophysics, by Carroll & Ostlie (2nd edition; Cambridge University Press, 2017). This book introduces more empirical knowledge (and jargon) assumed known in KW, and is used for the UofT undergraduate courses AST 221 and AST 320 (I'll try to refer to relevant notes from those classes; have a look at the parts on stars from the AST 320 notes and mini problem sets).
Evaluation
 Two problem sets (20% total), due two weeks after posting.
 Problem set 1 (pdf), due Feb 4.
 Problem set 2 (pdf), due Mar 4.
 Presentations (20% total), one short ones explaning a specific concept (8 min., plus 7 min. discussion; see list of topics) and one long one on a more advanced topic (list of topics; 20 min., plus 10 min. discussion).
 Computer project (20%), using
mesa
to investigate a specific part of stellar evolution.  Final exam (40%; likely oral, depending on number of students).
2 Master equations
Equilibria; timescales; massradius and massluminosity relations; HertzbrungRussell diagram, common threads in stellar evolution, features in stellar evolution.
Mon 7 Jan
 Textbook
 KW 1–2.3, 3.1 [and AST320 notes 1]
 Exercises
 What massradius relation would you expect for a set of stars with the same central temperature? Would more massive stars have higher or lower central density and pressure?
 AST320 mini problem sets: I, II
 General qualifier questions: Galactic:13, Stars:10,11, Physics:8 (first part)
 Derive the Virial Theorem (assuming ideal gas), following hint below.
 What is a star?
 me: ball of gas
 wiki:Star: A star is a massive, luminous sphere of plasma held together by gravity.
 Implication: P high → T high → loose heat
Unless high P without high T → degeneracy (brown dwarf, degenerate helium core, white dwarf).
 Basics of evolution: contraction, heat up, fusion
When does it stop?
 Where do we see stars
Check wiki:HRD. Why are there limits?
 left: T_{c} very high → fusion (think of contraction gone too far)
 top: L too high, matter blown away (L_{Edd})
 right: T profile too steep → convection
 Virial theorem: E_{int}=½ E_{grav}; E_{tot}=E_{int}+E_{grav}=½ E_{grav}
 Roughly, one has E_{grav} ≈ GM^{2}/R, E_{int} ≈ N⟨e_{int}⟩ = (M/μm_{H})3/2 kT, hence kT ≈ GMμm_{H}/R.
 To derive it formally, multiply HE by r on both sides and integrate over sphere; use that for ideal gas U = 3/2 nkT = 3/2 P.
Wed 9 Jan
 Textbook
 KW 2.4, 3.3, 5.1.1, 5.1.2, 20.2.2 [AST 320 notes 2]
 Exercises
 AST320 mini problem sets: III, VI
 Calculate dynamical time for the Sun (⟨ρ⟩≈1 g/cm^{3}), a neutron star (∼10^{14} g/cm^{3}), and the Universe as a whole (∼1 m^{3}≈10^{30} g/cm^{3}).
 Structure equations: MC, HE, TE, EB
 mass conservation: ∂m/∂r = 4πr^{2}ρ
 Hydrostatic equilibrium: ∂P/∂r = gρ = Gmρ/r^{2}
 Thermal equilibrium: ∂T/∂r = (3/4ac)(κρ/T^{3})(ℓ/4πr^{2}); easiest to derive from general diffusion equation: \(j = \frac13 vl_{mfp} \nabla{}n\); for radiation, v=c, l_{mfp }= 1/σn = 1/κρ, n = aT^{4}; hence, F_{rad} = ℓ/4πr^{2} = c/3κρ ∂(aT^{4})/∂r, from which one can solve for ∂T/∂r.
 Energy balance: ∂ℓ/∂r = 4πr^{2}ρε
 To solve, needs
 equation of state P(ρ,T,abundances)
 opacity κ(ρ,T,abundances)
 energy generation rate ε(ρ,T,abundances)
 If T (dependence) known, can solve MC, HE separately.
 Independent variable r or m (or P, or …)
 Homology
If two stars have the same structure, i.e., m'(r')/M'=m(r)/M for all r'/R' = r/R, then:
 MC: ρ'(r')/ρ(r) = (M'/M)(R'/R)^{3}
 HE: P'(r')/P(r) = (M'/M)^{2}(R'/R)^{4}
 Luminosity of a star
 Simple estimates/scalings:
 MC: ρ ≈ M/R^{3}
 HE: P ≈ GM^{2}/R^{4}
 TE: L ≈ acRT^{4}/κρ
 Combining with ideal gas law P=(ρ/μm_{H})kT:
 MC+HE: kT ≈ GMμm_{H}/R
 MC+HE+TE: L ≈ acG^{4}m_{H}^{4 }μ^{4}M^{3}/κ
Note: what is radiated does not depend on how energy is generated; star has to provide the energy, whether by contraction or fusion
 Simple estimates/scalings:
 Real MR and ML relations
 ε steep function of T → M/R nearly constant. Reality R ∝ M^{0.8}.
 κ not constant (higher at low T, but convection moderates the effect) → L ∝ M^{4}
 contribution of degeneracy → steeper at low M
 contribution of radiation pressure → flatter at high M.
 Substantial difference: inert lump inside → shell source
Or with fusion: ``star inside a star''
 Timescales
 Dynamical: τ_{dyn} ≈ 1/(Gρ)^{1/2}
What if not in HE?
Equation of motion: ρ∂v/∂t = ρ∂^{2}r/∂t^{2} = ∇P + ρ∇Φ = ∂P/∂r ρg
 Pressure drops away? τ_{ff} ≈ (R/g)^{1/2} ≈ 1/(Gρ)^{1/2}
 Gravity drops away? τ_{expl} ≈ R(P/ρ)^{1/2} ≈ R/c_{s} ≈ 1/(Gρ)^{1/2}
 Thermal: t_{th} ≈ E_{th}/L (≈ GM^{2}/RL for whole star)
If not in TE over some distance d: F = (vl_{mfp}/3)∇U ≈ (vl_{mfp}/3)U/d (where vl_{mfp}=c/κρ for diffusion by radiation).
Hence, timescale τ_{adj} ≈ Ud^{3}/Fd^{2} ≈ 3Ud^{3}/vl_{mfp}Ud ≈ d^{2}/vl_{mfp} ≈ (l_{mfp}/v)(d/l_{mfp})^{2} (random walk: t_{step} N_{steps})
Timescale for radiative damping of pulsations? Higher order → smaller d → faster damping.
 Nuclear: t_{nuc} ≈ E_{nuc}/L
 Dynamical: τ_{dyn} ≈ 1/(Gρ)^{1/2}
3 Equation of state
Fermions and bosons; pressure and energy density; ideal gas; (completely and partially) degenerate gas; radiation pressure; Boltzman distribution; Saha equation.
Mon Jan 14
 Textbook
 25.3 (esp. 25.3.2), 19.1–19.4, 19.9 [AST 320 notes 4]
 Exercises
 Write your own polytrope integrator (you'll need it for the first problem set; I suggest using python; if you are clueless, have a look at my simple integrator for an isothermal atmosphere)
 Use it to calculate the radius of a star with a solar mass and with central density and pressure like the Sun, for n=1.5 and 3.
 AST320 mini problem sets: V
 General qualifier questions: Physics:6
 Stability
Upon compression, ρ∝R^{3}. Thus, for an adiabatic perturbation, P ∝ ρ^{γ }∝ R^{3γ} (with γ a suitable average over the star). To keep in HE, P should increase as R^{4} or faster, i.e., 3γ>4 or γ>4/3.
 Polytropes: P = Kρ^γ ≡ Kρ^{1+1/n}
For K, γ, n ≡ 1/(γ1) constant, can integrate HE+MC.
Examples:
 Constant density (incompressible fluid)
 Isothermal (part)
 Completely convective
 Degenerate (K fixed)
For given K, n, know ρ(r), P(r), E_{g}, etc.; see AST 320 notes \#4, esp. Table 4.1.
Wed Jan 16
 Textbook
 Discussed in class: KW 13–16 (AST 320 notes 3); not discussed but to be read: KW 4, 25.3, 33.2
 Exercises
 For classical particles, show that n(p)dp is a Maxwellian, and that one recovers the ideal gas law.
 For photons, μ=0. Show that U_{ν}dν equals the Planck functions, and that its integral equals aT^{4} (note: ∫_{0}^∞ dx x^{3}/(exp(x)1)=π^{4}/15)
 AST320 mini problem sets: IV, VII
 EOS: Pressure integral: P = (1/3) ∫_{p} n_{p}v_{p}pdp
 NR: v_{p}p→p^{2}/m=2e_{p} → P=(2/3)U → Virial Theorem: E_{g}=2E_{i}, E_{tot}=(1/2)E_{g}
 ER: v_{p}p→cp=e_{p} → P=(1/3)U → Virial Theorem: E_{g}=E_{i}, E_{tot}=0
 generally, n_{p}=n(e_{p})g(4πp^{2}/h^{3})dp, n(e_{p})=1/[exp((eμ)/kT) ± 1]
+1:Fermions; 1: bosons; μ: chemical potential; g: number of internal states (such as spin)
(For a nice description of the meaning of The elusive chemical potential μ, see Baierlein 2001AmJPh..69..423B.)  ignore ±1: classical particles → ideal gas law: P=(ρ/μm_{H})kT (μ here is mean molecular weight)
 photons: 1, μ=0 → get BB → P=(1/3)aT^{4}
 electrons: +1: completely degenerate → fill up to p_{F} = h(3n/4πg)^{1/3}
 NRCD: P=K_{1}(ρ/μ_{e}m_{H})^{5/3}, K_{1}=(3/4πg)^{2/3}(h^{2}/5m_{e})^{nil}≈2.34×10^{38} N m^{3}
 ERCD: P=K_{2}(ρ/μ_{e})^{4/3} K_{2}=(3/4πg)^{1/3}(hc/4)≈2.45×10^{26} N m^{2}
 Complications: molecular/atomic/nuclear dissociation, pair formation
 Combinations
 Simplest: whichever dominates, or at least add radiation
 Rough estimate everywhere (Paczynski 1983ApJ...267..315P:
 P = P_{e}+P_{i}+P_{rad} = P_{e}+(ρ/μ_{i}m_{H})kT +(1/3)aT^{4}
 P_{e} = (P^{2}_{e,ideal}+P^{2}_{e,cd})^{1/2}, P_{e,ideal}=(ρ/μ_{e}m_{H})kT
 P_{e,cd} = (P^{2}_{e,nrcd}+P^{2}_{e,ercd})^{1/2}
 EOS from lookup table
 for completely ionised gas: Helmholtz (Timmes+Swesty 2000ApJS..126..501T; includes pair formation, explanation in Timmes+Arnett 1999ApJS..125..277T)
 MESA
 Another way to think about ionisation, etc. (not discussed in class)
(A bit different from KW 14)
Consider a fixed volume V at a fixed temperature T (or, equivalently, constant ρ and T). In thermal equilibrium, systems go to their most probable state, i.e., one maximizes entropy, \(S=k \log Z\), where \(Z\) is the partition function, a sum over all possible states i, weighted by \(\exp(E_i/kT)\). Usually, one can split contributions, e.g., for noninteracting photons, ions, and electrons, one has \(Z=Z_{\gamma}\times{}Z_{e}\times{}Z_{i}\) (and thus \(S=k\sum\log Z\)).
In the volume, for one particle at some momentum \(p\), the number of phase space elements available is \((V/h^3)\times4\pi{}p^{2}dp\), with a probability \(\exp(\epsilon_p/kT)\). The total number of phase space elements is thus \(\sim{}(V/h^3)p_{th}^3\), where \(p_{th}\) is some typical momentum associated with the temperature. Doing the integral gives the Maxwellian and \(p_{th}=\sqrt{2\pi{}mkT}\). Maybe more insightful is follow Baierlein 2001AmJPh..69..423B and define a typical size, \(\lambda_{th}\equiv{}h/p_{th}\), the ``thermal De Broglie'' wavelength. Then, the number of possible states is simply \(V/\lambda_{th}^3\). For a set of N identical particles, the contribution to the partition function is thus \[Z_N=\frac{[g(V/\lambda_{th}^3)\exp(\epsilon/kT)]^N}{N!},\] where \(g\) is the number of internal states, the factorial \(N!\) ensures we do not overcount states where two particles are swapped, and \(\epsilon\) is an energy cost beyond thermal kinetic energy there may be for having this particle.
Let's apply this to pair creation, assuming some mix of photons, ions, electrons and electronpositron pairs. Assuming a dilute plasma, their contributions to \(Z\) can be split, i.e., \(Z=Z_{\gamma}\times{}Z_{e}\times{}Z_{i}\times{}Z_{\pm}\) (of course, the physical picture is that there is a formation rate from the interactions of two photons, balanced by an annihilation rate; for the statistics, we are only concerned about the final equilibrium ). Since the electrons and positrons are independent, \(Z_{\pm}=Z_{+}\times{}Z_{}\), with both given by the above equation with \(\epsilon=m_{e}c^2\), but with \(N_{+}=N_{}=N_{\pm}\). Hence, \(Z_{+}=Z_{}\), and to find the number of particles, we can just find the maximum of \(S_{+}=k\log{}Z_{+}\), i.e., \[\frac{\partial{}S_{+}}{\partial{}N_{+}} = \frac{\partial{}k\log Z_{+}}{\partial{}N_{+}} = \frac{\partial}{\partial{}N_{+}}kN_{+}\left[\log\left(g\frac{V}{\lambda_{th}^3}\right)\frac{m_{e}c^2}{kT} \log N_{+}1\right]=0,\] where we used that for large \(N\), \(N!=N\log N  N\). Solving this for \(N_{+}\), one finds \[N=g\frac{V}{\lambda_{th}^3}\exp(m_{e}c^2/kT).\] Equivalently, one has \(n\equiv{}N/V=g\exp(m_{e}c^2/kT)/\lambda_{th}^3\), which has the nice implication that for classical particles, the probability for one with given internal state to exist in a given volume element \(\lambda_{th}^3\) is simple \(\exp(\epsilon/kT)\). Thus, for this very small volume, the probability becomes significant for \(kT\approx{}m_{e}c^2\). But when does the number of pairs become significant on larger scales? One measure to use is when \(n_{\pm}=n_e\), i.e., when \(\exp(m_{e}c^2/kT)=_{}\lambda_{th}^3 n_e/g\). For electrons (\(m=m_e\)), one has \(\lambda_{th}=2.4\times10^{10}T_9^{1/2}\) cm, and \(n_e=\rho/\mu_e{}m_H=6\times10^{23}(\rho_2/\mu_e)\) cm^{3}, so it requires \(T_{9}\approx{}m_{e}c^{2}/k(11.7+\log{}gT_9^{1/2}/\rho_2)\approx0.6\), quite consistent with KW, Fig. 34.1.
One can treat ionisation similarly, writing \(Z_H=Z_{0}\times{}Z_{p}\times{}Z_{e}\). We need to use that \(N_{p}=N_{e}=N_{H}N_0\). Doing a similar derivations as above, one derives the Saha equation. Again, ionisation is well before \(kT\approx{}\chi\). One consequence of this, is that if one, e.g., wants to know the population in excited states in hydrogen, it is easier to do this relative to the ionised state (since by the time you can excite even to the first excited state with \(\epsilon_2=\chi_H(11/4)=10.2\) eV, hydrogen is mostly ionised). For given state \(s\), one thus writes \(n(H_0,s)/n_{p} = (g_s/g_{p}g_{e}n_{e}\lambda_{th}^3)\times\exp((\chi\epsilon_2)/kT)\).
Finally, back to the chemical potential \(\mu\) (and Baierlein 2001AmJPh..69..423B). In terms of above quantities, one finds \(\mu=\epsilon+kT\log(g\lambda_{th}^3/n)\), but \(\mu\) also enters all thermodynamic potentials (internal energy U, enthalpy H, Helmholtz free energy F, Gibbs free energy G), as an additional term \(\dots+\mu{}dN\), i.e., the energy required to add one particle. In particular, for constant T, V, Helmholtz is handiest: \(F(T,V,N)=PV+\sum_{i}\mu_{i}N_{i}\) (and \(dF=PdV+\sum_{i}\mu_{i}dN_i\)). For pair plasma, minimizing \(F\) for \(N_{+}=N_{}\) (holding \(T\), \(V\), other \(N\) constant), one requires \(\mu_{+}+\mu_{}=0\). With the above microscopic definition of \(\mu\), one recovers the solution. Similarly, for ionisation, \(\mu_0=\mu_p+\mu_e\). In general, for any reaction left↔right, one expects that in equilibrium, \(\sum_{\rm left}\mu=\sum_{\rm right}\mu\). (In that sense, the above are missing photons – but these have \(\mu_\gamma=0\).)
All the above was for classical particles, but the same holds for nonclassical ones (except of course that one cannot assume a Maxwellian once particles start to overlap, \(\lambda_{th}\approx{}d=n^{1/3}\)). For completely degenerate neutron gas, where \(\mu=\epsilon_F\), one now trivially finds that there will be a contribution of protons and electrons such that \(\mu_n=\mu_p+\mu_e\). (Here, there is no \(\mu_\nu\), since the neutrinos escape; for a hot protoneutron star, where the neutrino opacity is still high, one does need to include it.) Remember, however, that above we derive a final, equilibrium state. The process to get there can be slow – not all baryons are in the form of iron yet!
 Gravothermal specific heat: c* = c_{P}(1∇_{ad}/∇_{gr}) (KW 4)
If fusion does not balance luminosity, heat/cool at a rate ∂q/∂t = ε∂l/∂m. Define: ε_{g} = ∂q/∂t = c_{P}T [(1/T)∂T/∂t  (∇_{ad}/P)∂P/∂t]
Here, following KW, S4.1, we used: dq = du+P dv = … = c_{P}(dT(∂T/∂P)_{s}dP) = c_{P}T(dT/T∇_{ad}dP/P), with ∇_{ad}=(∂lnT/∂lnP)_{s}
Important for contraction, on timescale τ≈τ_{th}. For slower changes (τ≫τ_{th}), negligible effect; for fast ones (τ≪τ_{th}), changes will be nearly adiabatic.
 Homologous change for sphere in HE (KW 25.3.4)
For small size change, homology gives dρ/ρ = 3dr/r, dP/P = 4dr/r and thus dρ/ρ = ¾dP/P. For a general EOS ρ ∝ P^{α}T^{δ} and thus dρ/ρ = αdP/P  δdT/T. Combining the two, we find that to keep HE, the temperature change has to be related to the pressure change by dT/T = (α¾)/δ dP/P ≡ ∇_{gr} dP/P.
Where we defined a temperature gradient ∇_{gr}=(∂lnT/∂lnP)_{gr} that is required for the star to remain in hydrostatic equilibrium.
Hence, the size change causes the change in heat per unit mass to be related to the temperature change via dq = c_{P}dT(1∇_{ad}/∇_{gr}) ≡ c* dT, with c* = c_{P}(1∇_{ad}/∇_{gr}) Compare with c_{P} at constant P, or c_{V} at constant volume; here it is at "constant hydrostatic equilibrium".
 Ideal gas: α=δ=1, ∇_{ad}=2/5, ∇_{gr}=1/4, hence c*=c_{P}(18/5)<0: removing heat causes the temperature to increase.
 Degenerate gas: α→3/5, δ→0, c*>0: removing heat causes the temperature to decrease.
 Homologous change for shell (KW 33.2)
Shell of width D at r, with D≪r, mass m~ρr^{2}D. Suppose expands with dr=dD, then dρ/ρ=dD/D. If outside expands homologously, still have dP/P=4dr/r=4(dD/D)(D/r). Hence, ∇_{gr} = (αD/r¼)/(δD/r) = (αr/4D)/δ and c*=c_{P}(1∇_{ad}δ/(αr/4D)).
Ideal gas: α=δ=1, unstable if D/r<¼. (think limit D→0: upon expansion, P does not change, ρ decreases, hence T must increase to keep HE).
 Homologous change for sphere in HE (KW 25.3.4)
4 Heat loss
Radiative diffusion; conduction; opacity sources; Schwarzschild and Ledoux criteria; mixing length theory; convective flux; stellar context for convection; semiconvection.
Mon Jan 21
 Textbook
 Discussed in class: KW 5.1; to be discussed as short presentations (and to be read): KW 5.2, 17 (AST320 notes 5)
 Exercises
 Redo AST320 miniPS VI, and think through what changes if you assume Kramerslike opacities (see KW 17.2–3).
 General qualifier questions: Stars:4
Wed Jan 23
 Textbook
 KW 6.1–6.5, 7, 30.4 (AST 320 notes 6)
 Exercises
 Study KW 24 (AST 320 notes 7) on the Hayashi line.
 General qualifier questions: Stars:8
 Criterion for convection: (1/γ)dlnP/dr > dlnρ/dr
 Schwarzschild criterion
 Ignoring composition gradients → ∇_{ad}<∇_{rad},
where ∇_{ad}=(dlnT/dlnP_{ad}=11/γ and ∇_{rad}=(dlnT/dr)_{rad}/(dlnP/dr)  Ledoux criterion
 With composition gradients → ∇_{ad}<∇_{rad}f∇_{μ},
where ∇_{μ}=dlnμ/dlnP and f=(∂lnρ/∂lnμ)/(∂lnρ/∂lnT); f=1 for fullyionised ideal gas.
 Convective flux
Generally, one can write the flux as, \[ F_{\rm conv} = \rho \overline{v}_{\rm conv} \Delta q = \rho \overline{v}_{\rm conv} c_P \Delta T, \] where \(\overline{v}_{\rm conv}\) is a ``suitable average'' of the convective velocity.
In terms of the gradients, one finds \[ F_{\rm conv} = \rho \overline{v}_{\rm conv} c_P T \frac{\ell_{\rm mix}}{2H_P}\left(\nabla\nabla_{\rm ad}\right), \] where \(\ell_{\rm mix}\) is the mixing length, usually parametrized as a fraction of the scale height, i.e., \(\ell_{\rm mix}\equiv\alpha_{\rm mix}H_P\), with \(\alpha_{\rm mix}\) the mixing length parameter.
The estimate of \(\overline{v}_{\rm conv}\) is the tricky part. We follow the AST 320 notes and balance buoyancy (\(Vg\Delta\rho=\rho Vg\Delta T/T\)) and friction (\(A\rho v^2\)); evaluate velocity at \(l_{\rm mix}/2\); define \(V/A=\beta\ell_{\rm mix}\), where \(\beta\) is a shape factor; and find \[ v_{\rm conv}^2 = \frac{\beta g}{H_P}\frac{\ell_{\rm mix}^2}{2} \left(\nabla\nabla_{\rm ad}\right). \] This leads to a convective flux given by \[ F_{\rm conv} = \rho c_P T \alpha_{\rm mix}^2 \sqrt{\frac{\beta g H_P}{8}} \left(\nabla\nabla_{\rm ad}\right)^{3/2}. \] Fortunately, the difficulty does not matter much: in the interiors of stars, convection is so efficient that the final temperature gradient ends up being essentially the adiabatic one. This is why we can treat completely convection stars as constantentropy polytropes. But near the atmosphere, this is no longer true.
Mon Jan 28
 Textbook
 KW 10, 24 (AST 320 notes 9, 7): Hayashi line.
 Exercises
 Check you understand the qualitative shapes of protostellar tracks (AST 320 notes, Fig. 7.3)
For fun, you could also have a look at the paper where the Hayashi line turned out to be important: 2000ApJ...529..428V (and the acknowledgement of the referee).
Wed Jan 30
 Textbook
 KW 17 (again; AST 320 notes 5)
 Exercises
 check you could do AST320 mini problem sets VII yourself (was presented by Jen)
 AST320 mini problem sets: XI, questions 1 and 2 (think ahead for question 3).
 Low temperature opacities: see Bethany's notes.
 Scalings for conduction
Aarya gave a nice qualitative introduction (see also KW 17.6). Here, a somewhat more mathematical one.
Generally, the flux is \(F=\frac13vl\nabla{}U\). It can be separated in different components. For photons, we saw \(U=aT^4\), \(v=c\) and \(l=1/\sigma{}n\) and hence one has \(F=(4ac/3)(T^3/\sigma{}n)\nabla{}T\) (where usually we write \(\sigma{}n=\kappa\rho\), but it is easier not to do so here). Given the definition of conductivity through \(F=k\nabla{}T\), one infers an equivalent conductivity \(k_{\gamma}=(4ac/3)(T^3/\sigma{}n)\).
For particles, \(U=\frac32nk_{B}T\) and thus \(F=\frac13vln\frac32k_{B}\nabla{}T\). Again writing \(l=1/n\sigma\), one finds \(k=\frac13\frac23k_{B}(v/\sigma)\). For an ideal, completely ionised gas, \(v\propto{}T^{1/2}\) and \(\sigma\sim{}Z^{2}e^4/(kT)^{2}\propto1/T^2\). Hence, \(k\propto{}T^{5/2}\).
For degenerate material, we should consider ions and electrons separately. The ions still have very short meanfree path, so do not contribute much. For the electrons, only a small fraction \(kT/E_{F}\) near the Fermi surface carries any heat, i.e., \(U_{e}\sim{}n_{e}(kT/E_F)kT\), and thus \(\nabla{}U\sim{}n_{e}(k_{B}T/E_{F})\nabla{}T\). Furthermore, those electrons have velocity depending on density, not temperature. Their free path still is \(l=1/n_{i}\sigma\) (\(n_i\) the ion density), but now \(\sigma\sim{}Z^{2}e^4/E_{F}^{2}\propto{}1/E_{F}^2\), and thus \(k_{e}\propto{}(v/\sigma{}n_{i})n_{e}(k_{B}T/E_{F})\propto{}vE_{F}T\). For nonrelativistic electrons, \(v\propto{}\rho^{1/3}\) and \(E_{F}\propto\rho^{2/3}\), so \(k_{e}\propto\rho{}T\). For relativistic particles, \(v\to{}c\) and \(E_{F}\propto\rho^{1/3}\), so \(k_{e}\propto\rho^{1/3}T\).
Writing in terms of an equivalent opacity, \(\kappa=(4ac/3)(T^3/k\rho)\), one finds for the ionised ideal gas, the opacity for electrons scales as \(\kappa_{e}\propto{}T^{1/2}/\rho\), for nonrelativistic degenerate electrons, \(\kappa_{e}\propto{}T^2/\rho^2\), and for relativistic degenerate electrons, \(\kappa_{e}\propto{}T^2/\rho^{4/3}\). Note that the photon opacity should also be affected, since photons can only interact with electrons near the Fermi surface, so \(l_{\gamma}\sim1/\sigma{}n_{e}(kT/E_{F})\). Equivalently, one can write that the effective opacity scales as \(T/E_{F}\propto{}T/\rho^{2/3}\) (nonrelativistic) or \(T/\rho^{1/3}\) (relativistic). At high densities, however, electron conduction will still win because of its steeper dependence on ρ.
5 Energy production
Nuclear binding energy; Coulomb barrier; reaction channels (PP, CNO, He and beyond, D/Li burning, s/r/pprocesses); rates and neutrinos.
Mon Feb 4
 Textbook
 KW 18.5.1 (pp and CNO cycles). Slowness of pp compared to
Li+p and D+p due to weak reaction.
Wed Feb 6
 Textbook
 KW 18
 Exercises
 AST320 miniPS XI, XII
 Temperature dependence
Generally, we write the cross section \(\sigma(E)=(S(E)/E)\exp(b/\sqrt{E})\), and integrate over \(E\) to get \(\langle\sigma{}v\rangle\), i.e., \[ \langle\sigma{}v\rangle=\sqrt{\frac{8}{\pi\mu}}\left(\frac{1}{kT}\right)^{3/2} \int_E S(E)\exp(E/kTb/\sqrt{E}) dE \] We discussed how normally \(S(E)\) can be taken out of the integral and one finds the Gamov peak, with height \(\exp(3E_0/kT)\), with \(3E_0/kT=19.721(\mu/m_u)^{1/3}(Z_{a}Z_{b})^{2/3}T_{7}^{1/3}\).
As Ariel mentioned, resonances are important. The above holds if one's energy is in the far wing of a resonance, so that \(S(E)\) indeed varies slowly. But if the resonance is inside the Gamov peak, it can dominate the energy dependence. In that case, one can consider it as a delta function, and the reaction rate will scale just with \(\exp(E_{res}/kTb/\sqrt{E_{res}})\), i.e., the only temperaturedependent part comes from how many particles have the right energy. For this reason, the \(3\alpha\) reaction rate has a term with \(\exp(C/T)\) instead of \(\exp(C/T^{1/3})\).
6 Evolution of single stars
Mon Feb 25  Main Sequence
 Textbook
 KW 28.2, 30, AST 320 notes 11.
 Exercises
 Read also KW 28.1 and AST 320 notes 11; former helpful for the MESA problem set as well.
 Check that you understand the different ends of the main sequence for different masses listed in AST 320 notes 11, and how these relate to what one sees in the HRD (AST 320 notes Fig 10.2 is well worth studying in detail).
 Approach to the main sequence
Generally, contract until some fusion process can provide the luminosity radiated. On the way to the main sequence, D and Li are fused, but for most stars, the first fusion stage that can hold up the contraction for a little while is the first part of the CNO cycle, where C is turned into N (see Fig. 7.3 in the AST 320 notes). Only when the C is exhausted does the star contract further until either the pp chain or the full CNO cycle takes over.
 On the main sequence
Hydrogen converted to Helium. In lowmass stars, radiative core so centre exhausts first. In more massive stars, convective core exhausts in one go, though the convection zone slowly becomes smaller during the main sequence. In detail, this depends on how convection actually happens, i.e., on overshooting and semiconvection.
For both, the luminosity increases slightly. Qualitatively, one can understand this from the increase in mean molecular weight μ. Naively, one would expect a decrease in radius, but changes in stellar structure counteract this (i.e., the star does not change homologously). Only in the final stages does the radius descrease a little.
 End of the main sequence
The core contracts and a shell around it ignites. In general, if a stable core can be formed, it will become isothermal. But there is a maximum (see KW and AST 320 notes); beyond that the core has to contract and either ignite He fusion or become degenerate.
Wed Feb 27  Giant Branch and Helium Flash
 Textbook
 KWW 33.1–6 (KW 32), AST 320 notes 12, lowmass giants.
 Exercises
 Check that you understand the basic differences between fusion in mainsequence and giant stars, and in degenerate cores.
 Ensure you understand why for lowmass stars, the helium flash happens at a fixed core mass and luminosity, (nearly) independent of the stellar mass.
 Study the figures with evolutionary tracks and perhaps especially Fig. 12.4 in AST 320 notes.
 Giant stars
For a sufficiently dilute envelope (M small and/or R large), the properties of shell determined by the core only, as the envelope is all ``far away.'' In particular, kT≈GM_{c}μm_{H}/R_{c}.
As a consequence, if, e.g., the core contracts, T will go up and so will the luminosity, causing the envelope to expand: mirror principle.
 Helium flash
For lowmass stars, the degenerate helium core is at about the same temperature as the shell. Eventually, helium ignites, at a core mass of about 0.45 M_{o}, somewhat off centre. Since the core is degenerate, a thermonuclear runaway ensues, though it does not become dynamically unstable.
Mon Mar 4  Intermediate Mass Giants, Blue Loops
 Textbook
 KWW 31, AST 320 notes 12, intermediatemass giants.
 Exercises
 Check you understand what causes the first and second dredge up.
 To better understand the loops, read Lauterborn et al., 1971A&A....10...97L (for recent discussion on blue loops, see Walmswell et al. 2015MNRAS.447.2951W; for more general insights, Gautchy 2018arXiv181211864G).
 Study the figures with evolutionary tracks and perhaps especially Fig. 31.2 in KW.
Wed Mar 6  Thermal Pulses on the AGB
 Textbook
 KWW 33.3 (RGB bump), 34.134.4 (AGB thermal pulses), AST 320 notes 12, thermal pulses
 Exercises
 Check you understand the RGB bump and its relation to blue loops.
 General qualifier questions: (part of) Physics:8
Mon Mar 11  Asymptotic Giant Branch
 Textbook
 KWW 34.45 (coremass luminosity relation; nucleosynthesis)
 Exercises
 Read also KWW 34.7 and ensure you understand the general evolution on the AGB.
 General qualifier questions: Stars:1, 9; part of Stars:20.
 Eddington luminosity
I find it easiest to derive from force balance (which makes sense only for optically thin material above a star's photosphere): \[ F_{\rm grav} = \frac{GM}{R^2}m = F_{\rm rad} = \frac{L}{4\pi{}R^2}\sigma{}N \] where with \(m=\rho{}V\) and \(\sigma{}N=\sigma{}nV=\kappa\rho{}V\), one finds \[ L_{\rm Edd} = \frac{\pi{}cGM}{\kappa}. \]
 End of the AGB
Near the end of the AGB, the luminosity from the shell approaches the Eddington lumnisity relevant for electronscattering opacity. At the cool photosphere, the opacity generally is smaller except when it gets cold enough for dust grains to form. This becomes particularly easy after C has been dredged up. (And pulsations help too.)
Wed Mar 13  InitialFinal Mass Relation and AGB Mass Loss
 Textbook
 KWW 34.6 (AGB mass loss, white dwarf initialfinal mass relation)
 Exercises
 Read KWW 34.9 about postAGB evolution.
 General qualifier questions: Stars 20 (Ariel's talk for r process)
Mon Mar 18  Overall Evolution, Supernovae
 Textbook
 KWW 34.8, 35, 36.1, 36.3 (up to 36.3.4); AST 320 notes 13
 Exercises
 Study both interior (ρT) and exterior (TL) diagrammes in detail,
ensuring you understand the basics. (Further nice ones in the
first
MESA
paper: Paxton et al. 2011ApJS..192....3P.)  AST320 mini problem sets: XIII
 Study both interior (ρT) and exterior (TL) diagrammes in detail,
ensuring you understand the basics. (Further nice ones in the
first
7 Binary evolution
Most stars increase in radius as they evolve, often drastically. If in a binary, they may at some point overflow their Roche lobes, leading to mass transfer to the companion. If this is stable, mass transfer will be on the evolutionary timescale. If unstable, it can be on the dynamical or thermal timescale. Masses transfer ceases when the star stops trying to expand; in giants, this is when most of the envelope has been transferred, and the remainder becomes so tenuous that it shrinks. Thus, one generally is left with just the core of the star. This process, and variations on it, is responsible for most of the more interesting stars we observe. For a general review, see Section 3 in Van den Heuvel, 2009ASSL..359..125V.
Figure 1: Radius evolution of stars of various masses. Lines indicate properties, as labelled; the one unmarked dotted line between ‘helium core flash’ and ‘core helium ignition’ marks the division between those helium cores (at lower masses) which evolve to degeneracy if stripped of their envelope, and those (at higher masses) which ignite helium nondegenerately and become helium stars. From Webbink 2008ASSL..352..233W, his Fig.~1.
Figure 2: Core masses as a function of radius and mass. Those interior to the hydrogenburning shell are indicated with solid lines, and dashed lines those interior to the heliumburning shell. Solid lines intersecting the base of the giant branch (dashdotted curve) correspond to helium core masses of to 0.15, 0.25, 0.35, 0.5, 0.7, 1.0, 1.4, and \(2.0\,M_\odot\); those between helium ignition and the initial thermal pulse to 0.7, 1.0, 1.4, and \(2.0\,M_\odot\), and those beyond the initial thermal pulse to 0.7, 1.0, and \(1.4\,M_\odot\). Dashed lines between helium ignition and initial thermal pulse correspond to carbonoxygen core masses of 0.35, 0.5, 0.7, 1.0, and 1.4 M_{o}. Beyond the initial thermal pulse, helium and carbonoxygen core masses converge, with the second dredgeup phase reducing helium core masses above about \(0.8\,M_\odot\) to the carbonoxygen core. From Webbink 2008ASSL..352..233W, his Fig.~2.
Mass loss and tranfer
Consider a star that looses or transfers mass at some rate \(\dot M\).
 Effect on orbit
The angular momentum of an orbit is given by \(J=(M_1M_2/M)\sqrt{GMa}\), and thus, \[ \frac{\dot J}{J} = \frac{\dot M_1}{M_1}+\frac{\dot M_2}{M_2} \frac{1}{2}\frac{\dot M}{M}+\frac{1}{2}\frac{\dot a}{a} \] With this, we can now consider several cases.
 Conservative mass transfer
Consider mass transfer from star 2 to star 1. If no mass and angular momentum is lost, then \(\dot M_1=\dot M_2\), \(\dot M=0\), \(\dot J=0\). Thus, \[ \frac{\dot a}{a} = 2\frac{M_{2}M_{1}}{M_{1}M_2}\dot M_{2} = 2(q1)\frac{\dot M_{2}}{M_{2}}, \] where \(q=M_2/M_1\) is the mass ratio between the donor (star 2) and the accretor (star 1). For donors less massive than the accretor, the orbit expands upon mass transfer (remember that \(\dot M_{2}<0\)).
Looking at the Roche lobe for a less massive donor, for which \(R_{L}\approx0.46a(M_{2}/M)^{1/3}\), one finds \[ \frac{\dot R_{L}}{R_{L}} = \frac{\dot a}{a} +\frac{1}{3}\frac{\dot M_{2}}{M_{2}} = 2\left(q\frac{5}{6}\right)\frac{\dot M_{2}}{M_{2}}, \] showing that the Roche lobe, as expected, grows a little slower than the orbital separation. (An analysis valid for all \(q\) would use the approximation of Eggleton 1983ApJ...268..368E, \(R_{L}/a\simeq 0.46q^{2/3}/[0.6q^{2/3}+\ln(1+q^{1/3})]\).)
 Spherically symmetric wind
\(\dot M_{2}=\dot M\), \(\dot M_{1}=0\), \(\dot J=(\dot M_{2}/M_{2})(M_{1}/M)J\). Hence, \[ \frac{\dot a}{a} = 2\left(\frac{M_{1} \dot M}{M_{2}M}\frac{\dot M}{M_{2}}+\frac{\dot M}{2M}\right) =\frac{\dot M}{M}. \] Thus, for mass loss (\(\dot M<0\)), the orbit expands.
 Spherically reemitted wind
\(\dot M_{2}=\dot M\), \(\dot M_{1}=0\), \(\dot J=(\dot M_{2}/M_{1})(M_{2}/M)J\) (idea is that accretor cannot handle mass transferred to it and reemits it as a wind). Hence, \[ \frac{\dot a}{a} = 2\left(\frac{M_{2}\dot M}{M_{1}M}\frac{\dot M}{M_{2}}+\frac{\dot M}{2M}\right) =\frac{2q^22q}{1+q}\frac{\dot M}{M}. \] Hence, orbit expands for \(q<(1+\sqrt{17})/4=1.28\) (with again a somewhat lower value for increasing Rochelobe radius), i.e., it is less quickly unstable than for conservative mass transfer. For a more detailed analysis, see Soberman et al., 1997A&A...327..620S
 Conservative mass transfer
 Effect on stellar radius
If the mass is lost from the outside of a star, the star becomes initially smaller, but on a hydrodynamic timescale it will partially reexpand in responds to the decreased pressure. Which effect dominates depends on the internal structure of the star. Generally, for thermal envelopes, the stars shrinks inside its Roche lobe, reexpanding only on the thermal timescale, typically to nearly its original size (especially for giants). However, a complication for thermaltimescale mass transfer is that, if the secondary is substantially less massive, it cannot accrete sufficiently fast and will bloat itself. For massive stars, this leads to contact, and almost certainly further mass loss and/or a merger.
Completely convective stars, or stars with deep convective layers, however, increase in size upon mass loss. For completely convective stars, which are described well by polytropes with \(P=K\rho^\gamma\) with \(\gamma=\frac{5}{3}\) (and thus \(n=1.5\)), this follows immediately from the mass radius relation: \(R\propto{}M^{1/3}\) (true for constant \(K\), i.e., for constant entropy or completely degenerate, nonrelativistic gas). Comparing this to the change in Roche lobe for conservative mass transfer, one sees that stability requires that \[ 2\left(q\frac{5}{6}\right)<\frac{1}{3} \Leftrightarrow q<\frac{2}{3}\qquad\hbox{for}\quad n=1.5. \]
Figure 3: \(\zeta_{\rm L}\equiv\partial\ln R_{\rm L}/\partial\ln M_2\) as a function of mass ratio, with all mass transfer through a single channel: conservative (cons); isotropic wind from donor star (wind); isotropic reemission of matter, from vicinity of `accreting' star (isor). (Also shown is a ring formation, indicative of mass loss from an outer Lagrange point). From Soberman et al., 1997A&A...327..620S, their Fig.~4.