Download Dr. Amanda Karakas and Prof. John Lattanzio

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