Transcript Slide 1

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PRESUPERNOVA EVOLUTION AND
EXPLOSION OF MASSIVE STARS:
CHALLENGES OF THE NEXT CENTURY
Marco Limongi
INAF – Osservatorio Astronomico di Roma, ITALY
email: [email protected]
Alessandro Chieffi
INAF – Istituto di Astrofisica Spaziale e Fisica Cosmica, ITALY
email: [email protected]
WHY ARE MASSIVE STARS IMPORTANT IN THE GLOBAL
EVOLUTION OF OUR UNIVERSE?
Chemical Evolution of Galaxies
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 z>5
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 events
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OVERVIEW OF MASSIVE STARS EVOLUTION
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Grid of 15 stellar models: 11, 12, 13, 14, 15, 16, 17,
20, 25, 30, 35, 40, 60, 80 and 120 M
Initial Solar Composition (A&G89)
All models computed with the FRANEC (Frascati RAphson Newton
Evolutionary Code) release 5.050419
(Limongi & Chieffi 2006, ApJ, 647, 483)

Evolution followed from the Pre Main Sequence up to the beginning
of the core collapse

4 physical + N chemical equations (mixing+nuclear burning) fully
coupled and solved simultaneously (Henyey)

Nuclear network very extended
282 nuclear species (H to Mo) and ~ 3000 processes (Fully Automated)
(NO Quasi (QSE) or Full Nuclear Statistical Equilibrium (NSE) approximation)

Convective Core Overshooting: d=0.2 Hp

Mass Loss: Vink et al. (2000,2001) (Teff>12000 K), De Jager (1988)
(Teff<12000 K), Nugis & Lamers (2000) (Wolf-Rayet)/Langer 1898
(WNE/WCO)
Convective
Core
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CORE H BURNING
g
g
g
CNO
Cycle
g
g
g
Core H
Burning
Models
g
g
Mmin(O) = 14 M
t(O)/t(H burning): 0.15 (14 M ) – 0.79 (120 M)
MASS LOSS
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CORE He BURNING
3a+
12C(a,g)16O
g
g
He
Convective
Core
g
Core He
Burning
Models
g
g
g
g
g
H burning shell
H exhausted core
(He Core)
M ≤ 30 M  RSG
M ≥ 35 M  BSG
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H
Log (Teff)>4.0
H/He
He/CO
H
H<0.4 WNL
He(N) WNE
WCO
He C,O
M ≥ 40
30 M
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MASSIVE STARS: MASS LOSS DURING H-He BURNING
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30  M / M O  35
RSG  WNL
35  M / M O  40
RSG  WNL  WNE
40  M / M O  60
RSG  WNL  WNE  WCO
60  M / M O
WNL  WNE  WCO
WR : Log10(Teff) > 4.0
WNL: 10-5< Hsup <0.4 (H
burning, CNO, products)
WNE: Hsup<10-5 (No H)
O-Type: 60000 > T(K) > 33000
WN/WC: 0.1 < X(C)/X(N) < 10
(both H and He burning
products, N and C)
WC: X(C)/X(N) > 10 (He burning
products)
 M < 30 M‫ סּ‬explode as Red SuperGiant (RSG)
 M ≥ 30 M‫ סּ‬explode as Blue SuperGiant (BSG)
ADVANCED BURNING STAGES
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Neutrino losses play a dominant role in the evolution of a
massive star beyond core He burning (T>109 K)
H-burn. shell
g
He core
He-burn.
shell
g

g

g

g



g
g

CO core
g

g  e  e   e   e
Core burning
g
The Nuclear Luminosity (Lnuc)
closely follows the energy losses
Enuc
L
M
t nuc
t nuc  Enuc
M
L
Evolutionary times of the advanced burning
stages reduce dramatically
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MASSIVE STARS: LIFETIMES
t nuc  Enuc
M
L
Fuel
Tc (K)
rc (g/cm3) Enuc (erg/g)
H
4.1(7)
4.7
6.44(18)
2.2(7) yr
(Lg=5·1038 )
6.87(6) yr
(Ltot=5·1038 )
He
2.1 (8)
7.2(2)
8.70(17)
1.6(6) yr
(Lg=9·1038 )
5.27(5) yr
(Ltot=9·1038 )
C
8.3(8)
1.7(5)
4.00(17)
6.3(5) yr
(Lg=1·1039 )
6.27(3) yr
(Ltot=8·1040 )
Ne
1.6(9)
2.5(6)
1.10(17)
1.7(5) yr
(Lg=1·1039 )
190 days
(Ltot=2·1043 )
O
2.1(9)
5.8(6)
4.98(17)
7.9(5) yr
(Lg=1·1039 )
243 days
(Ltot=9·1043 )
Si
3.5(9)
3.7(7)
1.90(17)
3.0(5) yr
(Lg=1·1039 )
19 days
(Ltot=2·1045 )
Estimated
Lifetime (no )
Real Lifetime
ADVANCED BURNING STAGES
Four major burnings, i.e., carbon, neon, oxygen and silicon.
Central burning  formation of a convective core
Central exhaustion  shell burning  convective shell
Local exhaustion  shell burning shifts outward in mass
 convective shell
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ADVANCED BURNING STAGES
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The details of this behavior (number, timing, overlap of
convective shells) is mainly driven by the CO core mass and
by its chemical composition (12C, 16O)
CO core mass
Thermodynamic history
12C, 16O
Basic fuel for all the nuclear burning
stages after core He burning
At core He exhaustion both the mass and the composition of the CO
core scale with the initial mass
ADVANCED BURNING STAGES
...hence, the evolutionary behavior scales as well
In general, one to four carbon convective shells and one
to three convective shell episodes for each of the neon,
oxygen and silicon burnings occur.
The number of C convective shells decreases as the mass
of the CO core increases (not the total mass!).
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PRESUPERNOVA STAR
The complex interplay among the shell nuclear burnings and
the timing of the convective zones determines in a direct way
the final distribution of the chemical composition...
16O,24Mg,
14N, 13C, 17O
28Si,29S,
30Si
28Si,32S,
36Ar,40Ca,
34S, 38Ar
12C, 16O
12C, 16O
s-
proc
20Ne,23Na,
56,57,58Fe,
52,53,54Cr,
55Mn,
59Co, 62Ni
NSE
24Mg,25Mg,
27Al,
s-proc
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PRESUPERNOVA STAR
...and the density structure of the star at the presupernova stage
The final Fe core Masses
range between:
MFe=1.20-1.45 M for
M ≤ 40 M
MFe=1.45-1.80 M for
M > 40 M
In general the higher is the mass of the CO core, the more
compact is the structure at the presupernova stage
THE SUPERNOVA PROBLEM
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The most recent and detailed simulations of core collapse SN
explosions show that:
the shock still stalls  No explosion is obtained
the energy of the explosion is a factor of 3 to 10
lower than usually observed
Work is underway by all the theoretical groups to better understand
the problem and we may expect progresses in the next future
The simulation of the explosion of the envelope is
mandatory to have information on:
the chemical yields (propagation of the shock wave 
compression and heating explosive nucleosynthesis)
the initial mass-remnant mass relation
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EXPLOSION AND FALLBACK
Different ways of
inducing the explosion
Shock Wave
Compression
and Heating
Induced
Expansion
and
Explosion
Initial
Remna
nt
• Piston (Woosley & Weaver)
• Thermal Bomb (Nomoto & Umeda)
• Kinetic Bomb (Chieffi & Limongi)
Matter
Falling
Back
Matter Ejected
into the ISM
Ekin1051 erg
Mass Cut
Initial
Remna
nt
Final
Remnant
Fe core
FB depends on the binding energy: the higher is the
initial mass the higher is the binding energy
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THE FINAL FATE OF A MASSIVE STAR
Z=Z
E=1051 erg
NL00
RSG
SNII
SNIb/c
WIND
WNE
WC/WO
Fallback
WNL
Black Hole
Neutron Star
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THE YIELDS OF MASSIVE STARS
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CONCLUSIONS. I
Stars with M<30 M explode as RSG
Stars with M≥30 M explode as BSG
The minimum masses for the formation of the
various kind of Wolf-Rayet stars are:
The final Fe core Masses range
between:
MFe=1.20-1.45 M for
MFe=1.45-1.80 M for
The limiting mass between SNII and SNIb/c is:
Salpeter IMF
WNL: 25-30 M
WNE: 30-35 M
WNC: 35-40 M
SNIb / c
 0.22
SNII
The limiting mass between NS and BH
formation is:
M ≤ 40 M
M > 40 M
30-35 M
SNII
SNIb/c
25-30 M
NS
M>35 M (SNIb/c) do not contribute to the intermediate and
heavy elements (large fallback)
BH
MAIN UNCERTAINTIES
PRESUPERNOVA EVOLUTION:
Mass Loss during Blue and Red supergiant phases, and
Wolf-Rayet stages
Treatment of Convection: extension of the convective
zones (overshooting, semiconvection), interaction
mixing-nuclear burning
12C(a,g)16O
cross section
Rotation
EXPLOSION:
Lack of an autoconsistent hydrodynamical model
(neutrino transport)
Induced explosion [Explosion energy (where and how),
time delay, fallback and mass cut (boundary conditions),
mixing (inner and outer borders), extra-fallback, Ye
variation, aspherical explosions]
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THE ROLE OF THE MASS LOSS FOR WNE/WCO
IN THE ADVANCED BURNING PHASES
Nugis & Lamers (2000) (NL00)
M  1011 ( L / L )1.29 Y 1.7 Z 0.5 M /yr

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Langer (1989) (LA89)
M  107 ( M / M )2.5 M / yr
Strong reduction of the He core during early core He burning
THE ROLE OF THE MASS LOSS FOR WNE/WCO
IN THE ADVANCED BURNING PHASES
LA89
NL00
LA89
NL00
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CONSEQUENCES ON THE FALLBACK
Final kinetic energy = 1 foe (1051 erg)
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THE FINAL FATE OF “LA89” MASSIVE STARS
Z=Z
E=1051 erg
LA89
SNII
SNIb/c
RSG
WNL
WNE
WC/WO
WIND
BH
Remnant Mass
NS
NS
25/34
THE YIELDS OF “LA89” MASSIVE STARS
NL00
LA89
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TREATMENT OF CONVECTION
Convection is, in general, a hydrodynamical multi-D phenomenon
 its inclusion in a hydrostatic 1-D stellar evolution code
consititutes a great source of uncertainty
Mixing-Length theory:
Extension of the convective zones (stability
criterion, overshooting, semiconvection)?
Temperature Gradient?
Interaction between nuclear burning and convective
mixing?
What about Mixing-Length theory for advanced burning stages
of massive stars?
dYi  Yi 
 
 Yi 
 Yi 
2




4

r
r






dt  t nuc  t conv  t nuc m 

Does it make sense?

2
Yi 
D
m 
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TREATMENT OF CONVECTION
PRODUCTION OF
60Fe
IN MASSIVE STARS:
He
60Fe
synthesize within the
He convective shell
X
Convection
T>4 108
produced
60Fe
22Ne,
a
Preserves 60Fe from
destruction
Brings new fuel (a,
22Ne)
He convective shell forms in a
zone with variable composition
X
M > 35 MO
M
M
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TREATMENT OF CONVECTION
Core Collapse and Bounce
Fe core

Shock Wave

Mass Fraction
THE MASS OF THE Fe CORE:
Si conv. shell
Si exhausted
Core
28Si
Final Fe Core Mass
Energy Losses
1 x 1051 erg/0.1M

Si conv. shell
Mass Fraction

M
Si exhausted
Core
28Si
Final Fe Core
Mass
M
UNCERTAINTY ABOUT
12C(a,g)16O
C0.2  X(12C)=0.2
C0.4  X(12C)=0.4
C0.4
C0.2
(Imbriani et al. ApJ 2001)
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30/34
THE ROLE OF ROTATION
Increasing rotation
OBLATENESS
Von Zeipel
Theorem
Cells of Meridional
Circulation
Frad  geff
GRATTONÖPIK CELL
Advection of Angular Momentum
Shear Instabilities:
- Mixing of chemical species
- Transport (diffusion)
of angular momentum
31/34
THE ROLE OF ROTATION
How include this multi-D phenomenon in a 1-D code?
Cylidrical Symmetry:
A(r , ,  )  A(r ,  )

A(r , )  A(r )  A(r ) P2 (cos  )
A(r )
A(r )
A(r , )
A(r ) 
 A sin  d
0
 sin  d
Average values over characteristic surfaces
Isobars/Equipotentials
A(r ,  )

0
1D problem
A(r )
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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
How much energy to inject and where:
Thermal Bomb (Internal Energy)
Kinetic Bomb (Initial Velocity)
Piston (Initial velocity and trajectory)
How much kinetic energy at infinity
[SN(~1051 erg)/HN(~1052 erg)]
Extension and timing (before/after fallback) of mixing
Efficiency of -process and changing of Ye
INDUCED EXPLOSION
Normal SN model (25M, E51=1)
Hypernova model with mixing-fallback
(25M, E51=10)
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Hypernova model without mixing-fallback
(25M, E51=10)
Hypernova model
(25M, E51=10, mixing-fallback, Ye)
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STRATEGIES FOR IMPROVEMENTS
Convection : hydrodynamical simulations in 3D  derive
simple prescriptions to be used in 1D hydrostatic models
(Arnett)
Rotation : implementation of 3D stellar models
12C(a,g)16O
: ask to nuclear physicists
Explosive Nucleosynthesis and Stellar Remnants : solve
the “Supernova Problem”  improve the treatment of
neutrino transport