MASSIVE STARS: PRESUPERNOVA EVOLUTION, EXPLOSION AND NUCLEOSYNTHESIS

1
MASSIVE STARS PRESUPERNOVA EVOLUTION, EXPLOSION
AND NUCLEOSYNTHESIS
Marco Limongi INAF Osservatorio Astronomico di
Roma, ITALY and Centre for Stellar and Planetary
Astrophysics Monash University AUSTRALIA Email
marco_at_oa-roma.inaf.it
2
What is a Massive star ?
It is a star that goes through all the
hydrostatic burnings in a quiescent way from H to
Si and eventually explodes as a core collapse
supernova
Mup
MPISN
lt Massive stars lt
gt120
8 - 10
3
Why are Massive stars important in the global
evolution of our Universe?
Light up regions of stellar birth ? induce star
formation
Production of most of the elements (those
necessary to life)
Mixing (winds and radiation) of the ISM
Production of neutron stars and black holes
Cosmology (PopIII)
Reionization of the Universe at zgt5
Massive Remnants (Black Holes) ? AGN progenitors
Pregalactic Chemical Enrichment
High Energy Astrophysics
Production of long-lived radioactive isotopes
(26Al, 56Co, 57Co, 44Ti, 60Fe)
GRB progenitors
The understanding of these stars, is crucial for
the interpretation of many astrophysical objects
4
BB Big Bang CR Cosmic Rays neut. n
induced reactions in SNII IMS Intermediate
Mass Stars SNII Core collapse
supernovae SNIa Termonuclear supernovae s-r
slow-rapid neutron captures
Le SNII contribuiscono in maniera rilevante
allevoluzione chimica della Galassia.
Responsabili per la nucleosintesi degli elementi
con 16ltAlt50 and 60ltAlt90
5
Computation of the Presupernova Evolution of
Massive Stars
1. Extended Network
Including a large number of isotopes and
reactions (captures of light partcles, e
captures, ß decays)
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Computation of the Presupernova Evolution of
Massive Stars
2. Strong coupling between physical and chemical
evolution

H/He burnings
Decoupled
Coupled
Adv. burnings
7
Computation of the Presupernova Evolution of
Massive Stars
3. Tratment of convection
- Time dependent convection
- Interaction between Mixing and Local Burning
D Diffusion Coefficient
8
Core H burning
Convective Core
g
g
g
CNO Cycle
g
g
g
g
g
Massive Stars powered by the CNO Cycle
The Convective Core shrinks in mass
9
CNO Cycle
12C 1H ? 13N g 13N ? 13C e
n 13C 1H ? 14N g 14N 1H ? 15O g
15O ? 15N e n 15N 1H ? 12C 4He
(99)

16O g
(1)
CN-Cycle
(T ? 3107 K)
16O 1H ? 17F g 17F ? 17O e
n 17O 1H ? 14N 4He
NO-Cycle
C? N? O?
CNO Processed Material
10
Ne-Na and Mg-Al Cycles
During Core H Burning the central temperature is
high enough (3-7107 K) that the Ne-Na and Mg-Al
cycles become efficient
Ne-Na Cycle
Mg-Al Cycle
20Ne 1H ? 21Na g 21Na ? 21Ne
e n 21Ne 1H ? 22Na g 22Na ?
22Ne e n 22Ne 1H ? 23Na g 23Na 1H
? 20Ne 4He
24Mg 1H ? 25Al g 25Al ? 25Mg
e n 25Mg 1H ? 26Al g 26Al ?
26Mg e n 26Mg 1H ? 27Al g 27Al 1H
? 24Mg 4He

• 21Na e 25Mg destroyed
• 22Ne slightly burnt
• 23Na e 26Mg increases
• 26Al (10-7) produced

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Evolutionary Properties of the Interior
t6.8 106 yr
12
Evolutionary Properties of the Surface
Core H Burning Models
Mmin(O) 14 M?
t(O)/t(H burning) 0.15 (14 M? ) 0.79 (120 M?)
13
Major Uncertainties in the computation of core H
burning models

• Extension of the Convective Core (Overshooting,
  Semiconvection)

• Mass Loss

Both influences the size of the He core that
drives the following evolution
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Core He burning
3a 12C(a,g)16O
4He 4He ? 8Be g 8Be ? 4He
4He 8Be 4He ? 12C g
g
He Convective Core
g
g
3 4He ? 12C g
g
g
g
g
g
H burning shell
H exhausted core (He Core)
The He convective core increases in mass
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Nucleosynthesis during Core He burning
3 4He ? 12C g 12C 4He ? 12O g 16O 4He ?
20Ne g 20Ne 4He ? 24Mg g
Chemical composition at core He exhaustion
mainly C/O
The C/O ratio is one the quantity that mainl
affects the advanced evolution of Massive Stars
(it determines the composition of the CO core)
C/O ratio depends on
1. Treatment of convection (late stages of core
He burning)
2. 12C(a,g)16O cross section
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Nucleosynthesis during Core He burning
14N, produced by H burning activates the sequence
of reactions
14N 4He ? 18F g 18F ? 18O e
n 18O 4He ? 22Ne g 22Ne 4He ? 25Mg n
For the CNO cycle
For e Solar composition
XCNO(iniziale) ? X14N
For a Solat composition at core H exhaustion
X(14N) ½ Z?
In general
The efficiency of the 14N reactions scales with
the metallicity
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Nucleosynthesis during Core He burning
14N ? 22Ne during the initial stages of core He
burning
He burning
H burning
CNO (1/2 Z)
14N (1/2 Z)
22Ne (Z)
During core He burning, 22Ne is reduced by a
factor of 2 by the nuclear reaction
22Ne 4He ? 25Mg n
Neutron Mass Fraction
s-process nucleosynthesis
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s-process during Core He burning
78Rb
79Rb
80Rb
81Rb
82Rb
83Rb
85Rb
84Rb
p
86Kr
87Kr
88Kr
77Kr
78Kr
79Kr
80Kr
81Kr
82Kr
83Kr
84Kr
85Br
86Br
87Br
b-
s
76Br
77Br
78Br
79Br
80Br
81Br
82Br
83Br
84Se
85Se
86Se
b-
75Se
76Se
77Se
78Se
79Se
80Se
81Se
82Se
83As
84As
85As
r
b-
74As
75As
76As
77As
80As
81As
78As
79As
b?
73Ge
74Ge
75Ge
76Ge
80Ge
77Ge
78Ge
79Ge
n,g
s,r
72Ga
73Ga
79Ga
76Ga
77Ga
78Ga
74Ga
75Ga
Both the neutron mass fraction and the seed
nuclei abundances scale with the metallicity
The abundance of the s-process nuclei scales with
the metallicity
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Evolutionary Properties of the Interior
WIND
t5.3 105 yr
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Evolutionary Properties of the Surface
Core He Burning Models
Core He Burning Models
M 30 M? ? RSG
M 35 M? ? BSG
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Major Uncertainties in the computation of core He
burning models

• Extension of the Convective Core (Overshooting,
  Semiconvection)

• Central 12C mass fraction (Treatment of
  Convection 12C(a,g)16O cross section)

• Mass Loss (determine which stars explode as RSG
  and which as BSG)

• 22Ne(a,n)25Mg (main neutron source for s-process
  nucleosynthesis)

All these uncertainties affect the size of the CO
core that drives the following evolution
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Advanced burning stages
Neutrino losses play a dominant role in the
evolution of a massive star beyond core He burning
At high temperature (Tgt109 K?0.08 MeV) neutrino
emission from pair production start to become
very efficient
g
n
n
He exhausted core (CO Core)
g
g
H burning shell
n
n
H exhausted core (He Core)
g
g
Core Burning
g
n
He burning shell
n
g
g
n
n
g
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Advanced burning stages
Evolutionary times of the advanced burning stages
reduce dramatically
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Evolutionary Properties of the Surface
M lt 30 M? ? Explode as RSG
M 30 M? ? Explode as BSG
After core He burning
At PreSN stage
Absolute Magnitude increases by 25
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Advanced Nuclear Burning Stages Core C burning
H burning shell
H
He
He burning shell
CO
T109 K
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Advanced Nuclear Burning Stages C burning
At high tempreatures a larger number of nuclear
reactions are activated
Heavy nuclei start to be produced
C-burning
Main Products of C burning
20Ne, 23Na, 24Mg, 27Al
Scondary Products of C burning
s-process nuclesynthesis
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Advanced Nuclear Burning Stages Core Ne burning
H burning shell
H
He burning shell
He
CO
NeO
C burning shell
T1.3109 K
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Advanced Nuclear Burning Stages Ne burning
Ne-burning
Main Products of Ne burning
16O, 24Mg, 28Si
Scondary Products of Ne burning
29Si, 30Si, 32S
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Advanced Nuclear Burning Stages Core O burning
H burning shell
H
He
He burning shell
CO
NeO
O
C burning shell
Ne burning shell
T2109 K
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Advanced Nuclear Burning Stages O burning
O-burning
Main Products of O burning
28Si (0.55)
32S (0.24)
Secondary Products of O burning
34S (0.07)
36Ar (0.02)
38Ar (0.10)
40Ca (0.01)
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Advanced Nuclear Burning Stages O burning
During core O burning weak interactions become
efficient
Most efficient processes
31S(b)31P
33S(e-,n)33P
30P(e-,n)30Si
37Ar(e-,n)37Cl
The electron fraction per nucleon
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Advanced Nuclear Burning Stages Core Si burning
H burning shell
H
He
He burning shell
CO
NeO
O
C burning shell
SiS
Ne burning shell
O burning shell
T2.5109 K
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Advanced Nuclear Burning Stages Si burning
At Oxygen exhaustion
Balance between forward and reverse (strong
interaction) reactions for increasing number of
processes
A measure of the degree of equilibrium reached by
a couple of forward and reverse processes
Non equilibrium
Full equilibrium
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Advanced Nuclear Burning Stages Si burning
At Oxygen exhaustion
At Si ignition
Sc
Si
Equilibrium
Equilibrium
Partial Eq.
Out of Eq.
Out of Equilibrium
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Advanced Nuclear Burning Stages Si burning

• 28Si is burnt through a sequence of (g,a)
  reactions

56Fe
A45
A44

• The two QSE clusters reajdust on the new
  equilibrium abundances of the light particles

28Si
24Mg

• The matter flows from the lower to the upper
  cluster through a sequence of non equilibirum
  reactions

20Ne
16O
12C
Equilibrium Clusters
Clusters di equilibrio
4He

• Ye is continuosuly decreased by the weak
  interactions (out of equilibrium)

56,57,58Fe, 52,53,54Cr, 55Mn, 59Co, 62Ni NSE
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Pre-SuperNova Stage
H burning shell
H
He
He burning shell
CO
NeO
O
C burning shell
SiS
Fe
Ne burning shell
O burning shell
T4.0109 K
Si burning shell
37
Evolutionary Properties of the Interior
H burning shell
Ne burning shell
He burning shell
O burning shell
C burning shell
Si burning shell
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Chemical Stratification at PreSN Stage
14N, 13C, 17O
14N, 13C, 17O
16O,24Mg, 28Si,29Si, 30Si
12C, 16O
28Si,32S, 36Ar,40Ca, 34S, 38Ar
12C, 16O s-proc
20Ne,23Na, 24Mg,25Mg, 27Al, s-proc
56,57,58Fe, 52,53,54Cr, 55Mn, 59Co, 62Ni NSE
Each zone keeps track of the various central or
shell burnings
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Main Properties of the PreSN Evolution
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Evolution of More Massive Stars Mass Loss
O-Type 60000 gt T(K) gt 33000
Wolf-Rayet Log10(Teff) gt 4.0

• WNL 10-5lt Hsup lt0.4 (H burning, CNO, products)
• WNE Hsuplt10-5 (No H)
• WN/WC 0.1 lt X(C)/X(N) lt 10 (both H and He
  burning products, N and C)
• WC X(C)/X(N) gt 10 (He burning products)

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Final Masses at the PreSN stage
HEAVY ELEMENTS
42
Major Uncertainties in the computation of the
advanced burning stages

• Treatment of Convection (interaction between
  mixing and local burning, stability criterion ?
  behavior of convective shells ? final M-R
  relation ? explosive nucleosynthesis)

• Computation of Nuclear Energy Generation (minimum
  size of nuclear network and coupling to physical
  equations, NSE/QSE approximations)

• Weak Interactions (determine Ye ? hydrostatic and
  explosive nucleosynthesis ? behavior of core
  collapse)

• Nuclear Cross Sections (nucleosynthesis of all
  the heavy elements)

• Partition Functions (NSE distribution)

• Neutrino Losses

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THE EVOLUTION UP TO THE IRON CORE COLLAPSE
The Iron Core is mainly composed by Iron Peak
Isotopes at NSE
The following evolution leads to the collapse of
the Iron Core
The Fe core begins to degenerate
The Fe core contracts to gain the energy
necessary against gravity
Tc 1010 K, rc 1010 K Pe 1028 dyne/cm2 Pi
21026 dyne/cm2 Prad 31025 dyne/cm2
The Chandrasekhar Mass MCh5.85(Ye)2 M? is
reached
T,r increase
A strong gravitational contraction begins
enuc lowers becaus the matter is at NSE
The Fermi energy increases?the electron captures
on both the free and bound protons incease as well
The gravitational collapse begins
The main source of pressure against gravity
(electron Pressure) lowers
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Fe Core
n
n
n-sphere
n
n
45
Strong Shock vs Weak Shock
A strong shock propagates. Matter is ejected.
A weak shock stalls. Matter falls back.
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Neutrino-driven explosions
Stalled Shock RS200-300 Km
Energy deposition behind the stalled shock wave
due to neutrino interactions
n heating
n diffusion
p,n
n cooling
n
n diffusion
Shock Wave reheated
e,e-
n,p
n
Explosion
n
Gain Radius RG100-150 Km
Neutrinosphere Rn50-700 Km
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Explosive Nucleosynthesis
Propatagiont of the shock wave through the
envelope
Compression and Heating
Explosive Nucleosynthesis
Explosion Mechanism Still Uncertain
The explosive nucleosynthesis calculations for
core collapse supernovae are still based on
explosions induced by injecting an arbitrary
amount of energy in a (also arbitrary) mass
location of the presupernova model and then
following the development of the blast wave by
means of an hydro code.

• Piston
• Thermal Bomb
• Kinetic Bomb

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Induced Explosion and Fallback
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Composition of the ejecta
The Iron Peak elements are those mostly affected
by the properties of the explosion, in particular
the amount of Fallback.
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The Final Fate of a Massive Star
51
Major Uncertainties in the simulation of the
explosion (remnant mass nucleosynyhesis)

• Prompt vs Delayed Explosion (this may alter both
  the M-R relation and Ye of the presupernova model)

• How to kick the blast wave
• Thermal Bomb Kinetic Bomb Piston
• Mass Location where the energy is injected

• How much energy to inject
• Thermal Bomb (Internal Energy)
• Kinetic Bomb (Initial Velocity)
• Piston (Initial velocity and trajectory)

• How much kinetic energy at infinity (typically
  1051 erg)

• Nuclear Cross Sections and Partition Functions

52
Chemical Enrichment due to Massive Stars
Different chemical composition of the ejecta for
different masses
53
Chemical Enrichment due to Massive Stars
Yields of Massive Stars used for the
interepretation of the chemical composition of
the Galaxy
We can have information on the contribution of
massive stars to the solar composition by looking
at the PFs of solar metallicity massive star
models.
ASSUMPTIONS
The average metallicity Z grows slowly and
continuously with respect to the evolutionary
timescales of the stars that contribute to the
environment enrichment
Most of the solar system distribution is the
result (as a first approximation) of the ejecta
of quasi solar-metallicity stars.
The PF of the chemical composition provided by a
generation of solar metallicity stars should be
flat
54
Chemical Enrichment due to Massive Stars
Yields averaged over a Salpeter IMF
Oxygen is produced predominantly by the
core-collapse supernovae and is also the most
abundant element produced by these stars
Use PF(O) to represents the overall increase of
the average metallicity and to verify if the
other nuclei follow or not its behavior
55
Chemical Enrichment due to Massive Stars
Elements above the compatibility range ? may
constitute a problem
Elements below the compatibility range ? produced
by other sources
Secondary Isotopes?
No room for other sources (AGB)
AGB
Type Ia
No room for AGB
n process. Other sources uncertain
Explosion?
56
Chemical Enrichment due to Massive Stars
Global Properties
IMF Salpeter
1 M?
Initial Composition (Mass Fraction)
Final Composition (Mass Fraction)
NO Dilution
Mrem0.186
X0.444 (f0.64) Y0.420 (f1.47) Z0.136
(f6.84)
X0.695 Y0.285 Z0.020
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Averaged Yields Relative Contributions
Stars with Mgt35 M? (SNIb/c) contribute for 20
at maximum (large fallback)
with few exceptions
(H,He burning)
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CONCLUSIONS

• Stars with Mlt30 M? explode as RSG Stars with M30
  M? explode as BSG

WNL 25-30 M?

• The minimum masses for the formation of the
  various kind of Wolf-Rayet stars are

WNE 30-35 M?
WNC 35-40 M?
MFe1.20-1.45 M? for M 40 M?

• The final Fe core Masses range between

MFe1.45-1.80 M? for M gt 40 M?
30-35 M?

• The limiting mass between SNII and SNIb/c is

SNII
SNIb/c
Salpeter IMF
25-30 M?

• The limiting mass between NS and BH formation is

NS
BH
(uncertainties on mass loss, simulated explosion,
etc.)
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CONCLUSIONS

• Massive Stars are responsible for producing
  elements with 4ltZlt38

• Assuming a Salpeted IMF the efficiency of
  enriching the ISM with heavy elements is

For each solar mass of gas returned to the ISM
H decreased by f0.64 He increased by
f1.47 Metals increased by f6.84

• SNIb/c contribute for 20 to the majority of the
  elements (large fallback)

Depends on Simulated expl. Mass Loss Binary
Systems ....... .......

• SNIb/c contribute for 40 to the elements
  produced by H and He burning that survive to
  fallback

Pre/Post SN models and explosive yields available
at http//www.mporzio.astro.it/limongi