Download FINAL STAGES OF MASSIVE STARS. SN EXPLOSION AND

FINAL STAGES OF MASSIVE STARS.
SN EXPLOSION AND EXPLOSIVE
NUCLEOSYNTHESIS
Marco Limongi
INAF – Osservatorio Astronomico di Roma, ITALY
Email: [email protected]
Why are Massive stars important in the global evolution of our
Universe?
Light up regions of stellar birth Æ 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) Æ 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 evidences
Log Mass Fraction
2
1
0
-1
-2
-3
-4
-5
-6
-7
-8
-9
-10
-11
-12
BB
Novae
SNIa
0
20
40
60
80
100
120
CR
IMS
s-r
140
neut.
SNII
160
180
200
Atomic Weight
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
Massive Stars contribute significantly to the chemical
evolution of the Galaxy through the production of elements
with 16<A<50 and 60<A<90
Massive Star:
It is a star that goes through all the nuclear burning stages
in a quiescent way (i.e., non degenerate environment) and
eventually explodes as a core collapse supernova
~8 M
AGB
~140 M
Massive Stars
PISN
Final Stages of Massive Stars:
All the hydrostatic nuclear burning stages that follow the
core He exhaustion and lead the star to the explosion
The various nuclear burnings involve nuclei of increasing Z
High Temperatures to overcome the Coulomb Barrier
Neutrino Losses from Pair Annihilation
At high temperatures (T>109 KÆ~0.08 MeV)
γ ↔ e+ + e− → ν e + ν e
The neutrinos exit the star @
the speed of the light while
positrons, electrons and γ’s
remain trapped
γ
γ
ν
ν
γ
ν
ν
From Carbon burning on
these losses greatly
dominate over radiative
diffusion and convection
γ
γ
γ
ν
γ
ν
ν
γ
ν
γ
Lifetimes of the Advanced Stages
Each burning stage gives about the same Enuc
Enuc
L≅
⋅M
t nuc
t nuc ≅ Enuc
M
L
Evolutionary times of the advanced burning
stages reduce dramatically
Lifetimes of the Advanced Stages
Fuel
Tc
ρc
H
3.7(7)
7.2
5.93(6) yr
He
1.5(8)
4.7(2)
6.8(5) yr
C
7.2(8)
1.2(5)
1.0(6)5.0(7)
4.0(7)1.0(9)
9.7(2) yr
Ne
1.2(9)
2.1(6)
7.0(9)
2.2(9)
280 days
O
1.8(9)
4.0(6)
5.0(10)
5.9(11)
4.0(10)
120 days
Si
3.1(9)
7.5(7)
1.1(13)
1.0(12)
7 days
Lnuc
Lν
Time
Synthesis of elements beyond O
At high temperatures light particles, released by key “ignition”
reactions can be captured by almost all the isotopes
i + j → k + light particle
These reactions produce
other light particles
(α,n)
(p,n)
A huge number of nuclear
reactions are activated
(p,γ)
(α,p)
β−,ε+
(n,γ)
(γ,n)
Heavy nuclei start to be
produced
(γ,p)
(p,α)
(γ,α)
(n,α)
(α,γ)
β+,ε−
(n,p)
Computation of the Presupernova Evolution of
Massive Stars
Zn
61
Cu
60
Ni
59
Co
58
Fe
57
60
57
54
52
Including a large number of
isotopes and reactions
(captures of light partcles,
e± captures, β± decays)
35
33
S
P
31
Si
30
P
30
Si
29
Al
28
29
Si
28
Al
27
27
He
25
1
H
2
H
3
H
Al
26
Cl
32
31
4
Ni
57
Co
56
Fe
55
32
Si
31
Sc
42
Ca
41
41
40
Cl
S
Ti
45
Sc
44
Ca
43
44
35
P
V
46
45
Ar
33
Co
55
Fe
54
53
Cu
59
Ni
58
Co
57
Fe
56
62
Zn
63
Cu
61
Ni
60
Co
59
Fe
58
Zn
64
Cu
62
Ni
61
Co
60
Fe
59
Zn
65
Cu
63
Ni
62
Co
61
Fe
60
Zn
66
Cu
64
Ni
63
Co
62
Fe
61
Zn
67
Cu
65
Ni
64
Zn
68
Cu
66
Cu
67
Ni
65
Ni
Co
Fe
Mn Mn Mn Mn Mn Mn Mn
48
36
34
Fe
51
37
He
58
56
1. Extended Network
3
Cu
Zn
K
38
Ar
37
Cl
36
S
34
P
33
Si
32
K
39
Ar
38
Cl
37
S
35
P
34
Si
33
S
K
Sc
43
Ca
42
40
Ar
39
Cl
38
36
S
K
Ar
41
40
K
Ar
Cr
V
47
Ti
46
Sc
45
Ca
44
42
41
49
Cr
V
48
Ti
47
Sc
46
Ca
45
52
50
Cr
V
49
Ti
48
Sc
47
Ca
46
53
51
Cr
V
50
Ti
49
Sc
48
Ca
47
54
52
Cr
V
51
Ti
50
Sc
49
Ca
48
55
53
Cr
V
52
Ti
51
Sc
50
Ca
49
54
56
Cr
55
57
Cr
V
Ti
Sc
Ca
K
Ar
Cl
37
S
P
Si
(α,n)
Al
(α,γ)
Mg 24Mg 25Mg 26Mg 27Mg
23
n
21
Na
22
Ne
21
20
F
18
O
17
16
17
15
10
7
Be
8
Li
7
6
Be
9
B
Be
O
16
13
N
14
N
15
N
12
C
13
C
14
C
11
10
F
19
O
18
Na
23
Ne
22
F
20
O
19
Na
24
Ne
23
Na
(p,n)
β−,ε+
Ne
(p,γ)
F
O
(n,γ)
(γ,n)
N
B
Be
(α,p)
(γ,p)
(p,α)
Li
(γ,α)
(n,α)
β+,ε−
(n,p)
Zn
Cu
Computation of the Presupernova Evolution of
Massive Stars
2. Strong coupling between physical and chemical evolution:
∂P
Gm
=−
∂m
4πr 4
∂r
1
=
∂m 4πr 2 ρ (P, T , Yi )
∂Yi
= ∑ ci ( j)λ jYj + ∑ ci ( j, k )ρN A < σv > j ,k YjYk
∂t
j
j ,k
∂L
= ε nuc ( P, T , Yi ) + εν ( P, T , Yi ) + ε grav (P, T , Yi )
∂m
∂T
GmT
=−
∇(P, T , Yi )
∂m
4πr 2 P
+
+ ∑ ci ( j, k , l ) ρ 2 N A < σv > j ,k ,l YjYkYl
2
j ,k ,l
i = 1,........,N
ρ = ρ ( P, T ) ; ε nuc = ε nuc ( P, T ) ; εν = εν ( P, T ) ;
H/He burnings:
ε grav = ε grav ( P, T ) ; ∇ = ∇( P, T )
Adv. burnings:
ε nuc = ε nuc ( P, T , Yi )
Decoupled
Coupled
Computation of the Presupernova Evolution of
Massive Stars
3. Tratment of convection:
- Time dependent convection
τ mix ≈ Δt
- Interaction between Mixing and Local Burning
∂P
∂m
∂r
∂m
∂L
∂m
∂T
∂m
∂Yi
∂t
τ mix ≈ τ nuc
Gm
4π r 4
1
=
4π r 2 ρ
=−
= ε nuc + εν + ε grav
GmT
∇
2
4π r P
⎛ ∂Y ⎞
⎛ ∂Y ⎞
=⎜ i ⎟ +⎜ i ⎟
⎝ ∂t ⎠ nuc ⎝ ∂t ⎠ conv
=−
∂ ⎛
∂Yi ⎞
⎛ ∂Yi ⎞
2
4
π
ρ
=
r
D
⎜
⎟
⎜
⎟
∂m ⎠
⎝ ∂t ⎠ conv ∂m ⎝
D = Diffusion Coefficient
Advanced Nuclear Burning Stages: C burning
At core He exhaustion the most abundant isotopes are
C-burning T ~ 109 K
12C,16O
Main Products of C burning
20Ne, 23Na, 24Mg, 27Al
Scondary Products of C
burning
22
Ne(α , n) 25Mg
s-process nuclesynthesis
Advanced Nuclear Burning Stages: Ne burning
At core C exhaustion the most abundant isotopes are
16O,20Ne
Ne-burning T ~ 1.3 ⋅109 K
Main Products of Ne burning
16O, 24Mg, 28Si
Scondary Products
of Ne burning
29Si, 30Si
Advanced Nuclear Burning Stages: O burning
O-burning
T ~ 2 ⋅109 K
Main Products of
O burning
28Si
(~0.55)
32S
(~0.24)
Secondary Products of
O burning
34S, 35Cl, 36,38Ar,
39K, 40,42Ca
Advanced Nuclear Burning Stages: O burning
Proton Number (Z)
During core O burning weak interactions become efficient
40Ca
41Ca
42Ca
43Ca
37K
38K
39K
40K
41K
42K
35Ar
36Ar
37Ar
38Ar
39Ar
40Ar
41Ar
33Cl
34Cl
35Cl
36Cl
37Cl
38Cl
31S
32S
33S
34S
35S
36S
37S
29P
30P
31P
32P
33P
34P
27Si
28Si
29Si
30Si
31Si
32Si
33Si
26Al
27Al
28Al
44Ca
Neutron Number (N)
Most efficient processes:
31S(β+)31P
33S(e-,ν)33P
30P(e-,ν)30Si
The electron fraction per nucleon
37Ar(e-,ν)37Cl
Zi
Ye = ∑ X i < 0.5
i Ai
Advanced Nuclear Burning Stages: Si burning
At Oxygen
exhaustion
T ~ 2.5 ⋅109 K
j+l
i+k
Balance between forward
and reverse (strong
interaction) reactions for
increasing number of
processes
rik = rjl
It is not possible to identify the most efficient sequence of
reactions anymore
A measure of the degree of equilibrium reached
by a couple of forward and reverse processes
φ (ij ) =
rik − rjl
φ (ij ) → 1
max( rik , rjl )
φ (ij ) → 0
Non equilibrium
Full equilibrium
Advanced Nuclear Burning Stages: Si burning
At Oxygen exhaustion
T ~ 2.5 ⋅109 K
φ (ij ) < 0.1
Sc
At Si ignition
At Si ignition
(panel a + panel b)
T ~ 3.5 ⋅109 K
T ~ 3.5 ⋅109 K
φ (ij ) < 0.01
Si
A=44
Equilibrium
Equilibrium
A=45
56Fe
28Si
0.01 < φ (ij ) < 0.1
φ (ij ) > 0.1
Partial Eq.
φ (ij ) < 0.1
Eq. Clusters
φ (ij ) > 0.1
Out of Equilibrium
Out of Eq.
In each equilibrium cluster the isotopes with the higher
binding energy are the most abundant ones
Advanced Nuclear Burning Stages: Si burning
T ~ 3.5 ⋅109 K
56
φ (ij) < 0.1
1.
Fe
A=45
A=44
28
16
12
Si
4
Mg
Ne
3. The matter flows
from the lower to the
upper cluster through
a sequence of non
equilibirum reactions
O
C
Equilibrium
Clusters
Clusters di equilibrio
He
is burnt through a
sequence of (γ,α) reactions
2. The two QSE clusters
reajdust on the new
equilibrium abundances of
the light particles
24
20
28Si
43
Ca ( n, γ ) 44 Ca
42
Ca (α , p) 45 Sc
44
Sc( n, p) 44 Ca
42
Ca (α , γ ) 46 Ti
42
Ca (α , n) 45 Ti
44
Ti (n, γ ) 45 Ti
44
Sc( p, γ ) 45 Ti
41
43
Ca (α , n) 46 Ti
44
41
K (α , p ) 44 Ca
4. Ye is continuosuly
decreased by the
weak interactions
(out of equilibrium)
Ca (α , γ ) 45 Ti
Sc(n, γ ) 45 Sc
56,57,58Fe, 52,53,54Cr, 55Mn,
59Co, 62Ni
NSE
Synthesis of the Elements
Fuel
Main Prod.
Sec. Prod.
ELEMENTS
NUMBER
1H
4He
13C, 14N, 17O
4,7
4He
12C, 16O
18O, 22Ne,
6,8
s-
proc.
12C
20Ne, 23Na,
25Mg,
s-proc.
10,11,12,13
24Mg, 27Al
20Ne
16O, 24Mg
29Si, 30Si
14
16O
28Si, 32S
Cl, Ar, K, Ca
14,16,17,18,
19,20
28Si
54Fe, 56Fe,
Ti, V, Cr, Mn,
Co, Ni
22,23,24,25,
26,27,28
55Fe
Convective History of a 20 M
H burning
shell
He burning
shell
Convective History of a 20 M
H burning
shell
He burning
shell
C burning
shell
C Convective
Core
C Convective
shells
Convective History of a 20 M
H burning
shell
He burning
shell
Ne Convective
Core
C burning
shell
Ne burning
shell
Convective History of a 20 M
H burning
shell
He burning
shell
C Convective
O
shells
C burning
shell
Ne burning
shell
O Convective
Core
O burning
shell
Convective History of a 20 M
H burning
shell
He burning
shell
Si Convective
shells
Si Convective
Core
C burning
shell
Ne burning
shell
O burning
shell
Si burning
shell
Convective History of a 20 M
H burning
shell
He burning
shell
Si burning
shell
C burning
shell
Ne burning
shell
O burning
shell
Pre-SuperNova Stage
H
He
CO
NeO
O
SiS
Fe
H burning
shell
He burning
shell
C burning
shell
Ne burning
shell
O burning
shell
T~4.0×109 K
Si burning
shell
Chemical Stratification at PreSN Stage
16O,24Mg,
14N, 13C, 17O
28Si,29Si,
30Si
12C, 16O
28Si,32S,
36Ar,40Ca,
12C, 16O
34S, 38Ar
s-proc
20Ne,23Na,
56,57,58Fe,
52,53,54Cr,
55Mn,
24Mg,25Mg,
27Al,
s-proc
59Co,
62Ni
NSE
Each zone keeps track of the various central or shell burnings
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 contracts to
gain the energy
necessary against gravity
T,ρ increase
No new fuel is found
At T ≥ 6 1010 K the
black body radiation
begins to disintegrate
the nuclei into free
nucleons
The Fe core begins to
degenerate
The Iron Core Mass
exceeds the
Chandrasekhar Mass
MCh=5.85×(Ye)2 M
A strong gravitational
contraction begins
The Fermi energy increasesÆthe
electron captures on both the free
and bound protons incease as well
The main source of pressure
against gravity (electron Pressure)
lowers
Tc ~ 1010 K, ρc ~ 1010 K
Pe ~ 1028 dyne/cm2
Pi ~ 2×1026 dyne/cm2
Prad ~ 3×1025 dyne/cm2
The Fe core
contracts rapidly
T,ρ increase
The gravitational
collapse begins
ρ ≈ 3 ⋅1012 g/cm 3
ρ ≈ 3 ⋅1011 g/cm 3
Fe Core
e− + p →ν e + n
ν
ν
ν
ν
ν
ν
ρ ≈ 10 g/cm
14
Fe Core
ν diffusion
ν
ν-sphere
ν
Neutrino Trapping
3
ν
Shock
wave
ν
ν
ν
Core Bounce and
Rebounce
Energy Losses
1 x 1051 erg/0.1M
Stalled Shock
ν
ν
ν
ν
“Prompt”shocks eventually stall!
Strong Shock vs Weak Shock
A strong shock propagates.
Matter is ejected.
A weak shock stalls.
Matter falls back.
Neutrino-driven explosions
Stalled Shock
RS=200-300 Km
ν heating
ν cooling
ν
ν diffusion
ν
p,n
ν e ,ν e
Energy deposition behind
the stalled shock wave
due to neutrino
interactions:
−
e
+
e
ν +n→ p+e
ν + p →n+e
Shock Wave reheated
e-,e+
n,p
Explosion!
ν
Gain Radius
RG=100-150 Km
Neutrinosphere
Rν=50-700 Km
Explosive Nucleosynthesis
Propagation 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
Induced Explosion and Fallback
Induced
Shock
Compression
and Heating
Induced
Expansion
and
Explosion
Initial
Remnant
Injected
Energy
Matter Ejected
into the ISM
Ekin≈1051 erg
Matter
Falling
Back
Mass Cut
Initial
Remnant
Final
Remnant
Explosion History of a 15 M
Since nuclear reactions are
very temperature sensitive,
this cause nucleosynthesis to
occur within few seconds
that might otherwise have
taken days or years in the
presupernova evolution.
Basic Properties of the Explosion
• Behind the shock, the pressure is dominated by radiation
• The shock propagates adiabatically
Shock
Fe
core
T1
r1
T2
r2
r
4
3
4
E = πaRPN Tpeak
3
The peak temperature does not depend on the stellar structure
Properties of the matter at high temperature
• If T>5 109 K all the forward and the reverse strong reactions come to
an equilibrium
• The abundances of the various nuclei are determined by the
conditions that each isotope is in equilibrium with the “sea” of
light particles:
A( N , Z ) ↔ Zp + Nn
⎛ 2πh
N + Z −1 ω ( N , Z )
2
Y ( N , Z ) = ( ρN A )
( N + Z ) ⎜⎜
N +Z
2
⎝ mH kT
3
2
⎞
⎟
⎟
⎠
3 ( N + Z −1)
2
NUCLEAR STATISTICAL EQUILIBRIUM
e
−
Q( N ,Z )
kT Y Z Y N
p n
Properties of the matter at high temperature
Yi NSE = f (T , ρ ,Ye )
T = 5 ⋅ 109 K ρ = 108 g/cm3
Ye=0.500, 56Ni
Ye=0.481, 54Fe
Ye=0.464, 56Fe
Ye=0.448, 58Fe
Since the matter exposed to the
explosion has Ye>0.49 (η<0.02)
Most abundant isotope
Elements also produced: Sc, Ti, Co, Ni, Zn
56Ni
Normal Freezout
α-rich Freezout
Properties of the matter at high temperature
• If T < 5 109 K not all the processes come to equilibrium.
The first processes that go out of the equilibrium are
those corresponding to A~44 that act as a bottleneck
A=44
A=45
56Fe
QUASI-EQUILIBRIUM (QSE)
28Si
Eq. Clusters
Yi QSE = f (T , ρ , Ye , 28Si)
Since the matter exposed to the explosion has A<44
The matter remains
partially locked as 28Si
Elements produced: V, Cr, Mn
Properties of the matter at high temperature
• If T < 4 109 K the processes corresponding to A~44 almost
inhibit the flux of the matter through the bottleneck
A=44
A=45
56Fe
TWO SEPARATE QSE CLUSTERS
28Si
Eq. Clusters
QSE
Ycluster
= f (T , ρ , Ye , ∑ ni )
The matter remains locked
into the lower cluster
Elements produced: Si, S, Ar,, K, Ca
Q
Properties of the matter at high temperature
• If T < 3.3 109 K the processes are far from the equilibrium
and nuclear processing occur through a well defined
sequence of nuclear reactions.
Elements preferrentially synthesized in these conditions
over the typical eplosion timescales:
2.1< T (K) < 3.3
T (K) < 2.1
Mg, Al, P, Cl
Ne, Na
• If T < 1.9 109 K no nuclear processing occur over the typical
explosion timescales.
By combining the properties of the matter at high temperature
and the basic properties of the explosion
Complete
Si
burning
T > 5 ⋅109 K
NSE
Incomplete
Si burning
T > 4 ⋅109 K
4
Explosive
O burning
T > 3.3 ⋅109 K
Eexpl = 1 foe
Explosive
Ne burning
T > 2.1⋅109 K
Explosive
C burning
T > 1.9 ⋅109 K
QSE
QSE
1 Cluster 2 Clusters
Sc
Ti
Fe
Co
Ni
Cr
V
Mn
3700
5000
Ne
Na
Mg
Al
P
Cl
Si
S
Ar
K
Ca
6400
RADIUS (Km)
11750
No Modification
⎛ 3Eexpl ⎞
⎟
Tmax = ⎜⎜
3 ⎟
⎝ 4πaRPN ⎠
1
13400
Role of the Progenitor Star
• Mass-Radius relation @ Presupernova Stage:
determines the amount of mass contained in each
volume Æ determines the amount of mass processed by
each explosive burning.
T > 5 ⋅109 K
NSE
Sc
Ti
Fe
Co
Ni
Incomplete
Si burning
T > 4 ⋅109 K
Explosive
O burning
T > 3.3 ⋅109 K
Explosive
Ne burning
T > 2.1⋅109 K
QSE
QSE
1 Cluster 2 Clusters
Si
S
Cr
Ar
V
K
Mn
Ca
INTERIOR MASS
Mg
Al
P
Cl
Explosive
C burning
T > 1.9 ⋅109 K
Ne
Na
No Modification
Complete
Si
burning
Role of the Progenitor Star
• The Ye profile at Presupernova Stage:
it is one of the quantities that determine the chemical
composition of the more internal zones that reach the
NSE/QSE stage
ρ=108 g/cm3
T=5·109 K
Ye=0.50 Æ
56Ni=0.63
Ye=0.49 Æ
54Fe=0.28
–
–
55Co=0.11
–
52Fe=0.07
–
57Ni=0.06
–
54Fe=0.05
56Ni=0.24
–
55Co=0.16
–
58Ni=0.11
–
57Ni=0.08
• The Chemical Composition at Presupernova Stage:
it determines the final composition of all the more
external regions undergoing explosive (in non NSE/QSE
regine)/hydrostatic burnings
The chemical composition of a massive star
after the Explosion
EXPLOSIVE BURNINGS
T > 5 ⋅109 K
Incomplete
Si burning
T > 4 ⋅109 K
Explosive
O burning
T > 3.3 ⋅109 K
QSE
QSE
Sc,Ti,Fe 1 Cluster 2 Clusters
Si,S,Ar
Co,Ni
Cr,V,Mn
K,Ca
Explosive
Ne burning
T > 2.1⋅109 K
Explosive
C burning
T > 1.9 ⋅109 K
NSE
CC = f (T , ρ , Ye )
Mg,Al,P,Cl
Ne,Na
CC = f (T , ρ , X i )
INTERIOR MASS
No Modification
Complete
Si
burning
Fallback and the Final Remnant
During the propagation of the shock wave through the mantle
some amount of matter may fall back onto the compact remnant
It depends on the binding energy of
the star and on the final kinetic
energy
Composition of the ejecta
The Iron Peak elements are those mostly affected by the properties
of the explosion, in particular the amount of Fallback.
The Final Fate of a Massive Star
Z=Z
E=1051 erg
NL00
oss
L
ass
M
No
WIND
SNII
SNIb/c
RSG
ss
Ma
e
or
ss
C
a
M
He
C
C
ss
a
He e M
r
Co
CO
WNL
WNE
R
WC/WO
nt
a
n
em
ss
a
M
Black Hole
Fe-Core Mass
Neutron Star
Fallback
al
F in
s
Mas
....this kind of pictures MUST be taken with caution
The impact of the boundary conditions:
Another possbile Final Fate of a Massive Star
The Role of Mass Loss
In stars with M>25 M the mass loss plays an important role
The Role of Mass Loss
In stars with M ≥ 40 M :
During core H burning the region of variable H is lost by stellar
wind Æ some isotopes are ejected by wind before their further
destruction (14N,23Na, 26Al)
The He convective core progressively receids in mass and leaves a
region of variable chemical composition that reflects the central
burning
A fraction of this region may be lost by stellar wind in the more
massive models where the mass loss is more efficient
After core He exhaustion a He convective shell forms in the region
with the He profile. The high temperatures induce the synethesis
of some nuclear species like, e.g., 20Ne, 24Mg in the more massive
models, and 60Fe
At the presupernova stage the chemical composition of the He
convective shell reflects both the central and shell burnings
In spite of the large remnant masses a large fraction of matter
processed by the He convective core is ejected in the ISM
Chemical Enrichment due to a single Massive Star
The Production Factors (PFs) provide information on the global
enrichment of the matter and its distribution
Mtot
Mtot
∫ X dm
i
PFi =
Mcut
Mtot
Initial
i
Mcut
∫X
dm
Solar Metallicity
Models
PFi =
∫ X i dm
Mcut
Mtot
Sun
X
i
∫ dm
Mcut
Chemical Enrichment due a generation of
Massive Stars
The integration of the yields provided by each star over an initial
mass function provide the chemical composition of the ejecta due
to a generation of massive stars
Yields averaged over a
Salpeter IMF
120
Yi
tot
= ∫ Yiφ (m)dm
11
φ (m) = km −α
α = 2.35
The Role of the More Massive Stars
Which is the contribution of stars with M ≥ 35 M?
Mass Loss Prevents Destruction
Large Fall Back
They produce:
~60% of the total C and N (mass loss)
~40% of the total Sc and s-process elements (mass loss)
No intermediate and iron peak elements (fallback)
Chemical Enrichment due to Massive Stars
Global Properties:
IMF: Salpeter
11M
M
Initial Composition
(Mass Fraction)
X=0.695
Y=0.285
Z=0.020
NO Dilution
Mrem=0.186
Final Composition
(Mass Fraction)
X=0.444 (f=0.64)
Y=0.420 (f=1.47)
Z=0.136 (f=6.84)
Chemical Enrichment due to Massive Stars
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 PFs of the chemical composition provided by a
generation of solar metallicity stars should be flat
Chemical Enrichment due to Massive Stars
No room for other
sources (AGB)
Remnant Masses?
Secondary
Isotopes?
ν process.
Other sources
uncertain
Type Ia
Explosion?
AGB?
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:
WNL: 25-30 M
WNE: 30-35 M
WNC: 35-40 M
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
SNIb / c
≅ 0.22
SNII
M ≤ 40 M
M > 40 M
30-35 M
SNII
SNIb/c
25-30 M
The limiting mass between NS and BH formation is:
(uncertainties on mass loss, simulated explosion, etc.)
NS
BH
CONCLUSIONS. I
Massive Stars are responsible for producing elements with 4<Z<38
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 f=0.64
He: increased by f=1.47
Metals: increased by f=6.84
M>35 M (SNIb/c) do not contribute to the intermediate mass
elements (large fallback)
M>35 M (SNIb/c) contribue for ~60% to the production of C and F,
and for ~60% to the production of the s-process elements (i.e.,
elements produced by H and He burning that survive to further
burning and fallback)
Pre/Post SN models and explosive yields available at
http://www.mporzio.astro.it/~limongi
MAIN UNCERTAINTIES AND OPEN PROBLEMS
H Burning:
Extension of the Convective Core
(Overshooting, Semiconvection)
Mass Loss
Both influence the size of the He core that
drives the following evolution
He Burning:
Extension of the Convective Core
(Overshooting, Semiconvection)
Central
12C
Convection +
mass fraction (Treatment of
12C(α,γ)16O
cross section)
Mass Loss (determine which stars explode
as RSG and which as BSG)
All these uncertainties affect the size and the
composition of the CO core that drive the
following evolution
MAIN UNCERTAINTIES AND OPEN PROBLEMS
Advanced Burnings:
Treatment of Convection (interaction between mixing and
local burning, stability criterion Æ behavior of convective
shells Æ final M-R relation Æ explosive nucleosynthesis)
Computation of Nuclear Energy Generation (minimum
size of nuclear network and coupling to physical equations,
NSE/QSE approximations)
Weak Interactions (determine Ye Æ hydrostatic and
explosive nucleosynthesis Æ behavior of core collapse)
Nuclear Cross Sections (nucleosynthesis of all the heavy
elements)
Partition Functions (NSE distribution)
Neutrino Losses
MAIN UNCERTAINTIES AND OPEN PROBLEMS
Explosion: Explosive Nucleosynthesis, Remnant Masses:
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
Boundary Conditions (Transmitting, Reflecting)
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
EVOLUTION AND EXPLOSION OF
POPIII (Z=0) MASSIVE STARS
Impact of initial metallicity on the evolution of a massive star
T of H burning:
Z=Z Æ CNO sustained Æ T ~ 4 107 K
Z=0
Æ Lack of CNO (PP-chain) Æ T ~ 108 K
H convective core:
Z=0
Æ lower opacity κ Æ lower ∇ rad
Æ energy production zone more
extended Æ lower flux
Smaller H convective cores
Smaller He core masses
Impact of initial metallicity on the evolution of a massive star
Z=0 Æ Low opacity Æ No expansion and cooling
Æ No convective envelope
Mass Loss would significantly affects the evolution of these
stars, but we do not have any information about this
phenomenon @ Z=0
Impact of initial metallicity on the evolution of a massive star
@ Z=0 the mixing of the H-rich mantle with the active He burning
region is quite COMMON because of the low entropy barrier that
exists at the H-He interface.
As a consequence, p and
12C
activate the CNO cycle Î large
production
14N
But also
production of Ne-Na-Mg-Al via p captures (Ne-Na, Mg-Al cycles)
Impact of initial metallicity on the evolution of a massive star
20 < M / M O < 50
Primary
14N
Primary
14N
production
production in the He convective shell due to the
injestion of H rich layers
Impact of initial metallicity on the evolution of a massive star
The basic evolutionary properties of the advanced phases after
core He burning mildly depend on the initial metallicity
Z=0
Z=Z
Impact of initial metallicity on the progenitor structure:
The Initial Mass – Remnant Mass Relation
Zero metallicity stars turn out to be more compact than the
solar metallcity ones
Ekin=1.2 foe
Zero metallicity stars leave larger remnants than
the solar metallicity ones
M > 25 M Æ
heavy element synthesis negligible
Impact of initial metallicity on the nucleosynthesis
Direct:
Initial Isotopic Distrubution (seeds)
CNO Æ
14N
Æ
22Ne
Æ neutrons
Electron fraction profile Ye
Odd – Even effect
Affects the explosive
burnings. NSE=f(ρ,T,Ye)
Indirect:
The initial metallicity influences the evolutionary properties
(convective cores, core masses....) Æ nucleosynthesis during the
various burning stages
Impact of initial metallicity on the chemical
yields of a massive stars
Yields averaged over a Salpeter IMF
φ (m) = km −α
Large Even-Odd effect
No production above Zn
α = 2.35
The Yields of the PopIII CC-SNe and the abundances in
extremely metal poor stars
Extremely Metal (Iron) Poor Stars = [Fe/H]<-3.0
⎛ FeStar
⎡ Fe ⎤
⎛ Fe ⎞
⎛ Fe ⎞
⎜
=
−
Log
Log
Log
=
⎜
⎟
⎜
⎟
⎢H⎥
⎜ Fe
⎣ ⎦
⎝ H ⎠Star
⎝ H ⎠SUN
⎝ SUN
⎡ Fe ⎤
Log ( FeStar ) = ⎢ ⎥ + Log ( FeSUN )
⎣H⎦
Fe EMPS < 10−6
⎞
⎛H
⎟ − Log ⎜ SUN
⎟
⎜H
⎠
⎝ Star
⎞
⎟
⎟
⎠
The Yields of the PopIII CC-SNe and the abundances in
extremely metal poor stars
Extremely Metal (Iron) poor stars probably formed in the
very early epochs of Galaxy formation by gas clouds
chemically enriched by the first stellar generations
(POPIII Core Collapse Supernovae)
Weather or not they are associated to single supernovae or
single burst events:
1. They provide very useful constraints to test presupernova
models, supernova explosions and nucleosynthesis theories
2. They can be used to infer the nature of the first
generations of stars and supernovae
Interpretation of “NORMAL” EMPS in terms of the ejecta of one or more
POPIII Core Collapse Supernovae
Limongi & Chieffi 2002, PASA, 19, 246
Chieffi & Limongi 2002, ApJ, 577, 281
The abundance pattern of the Normal EMP stars can be convincingly
explained in terms of the ejecta of one or more core collapse supernovae
HE0107-5240
• Log10(g) = 2.2
Log10(Teff)=3.707
[Fe/H]=-5.3
• C, N, Na enhancements by a factor of 104, 102.3 and 10 relative to Fe
• Mg, usually enhanced in EMPS, is at the level of Fe
[Fe/H]=-5.3
Table 1 Elemental abundances for HE010725240
Element
[X/Fe]
.............................................................................................................................................................................
Li
C
N
Na
Mg
Ca
Ti
Ni
Zn
Sr
Ba
Eu
,5.3
4.0
2.3
0.8
0.2
0.4
20.4
20.4
,2.7
,20.5
,0.8
,2.8
.............................................................................................................................................................................
Abundance ratios [X/Fe] of HE010725240 as derived from a high-resolution, high-S/N UVES
spectrum. In our analysis, we used a custom plane-parallel model atmosphere with the most
recent atomic and molecular opacity data. Typical errors in the logarithmic abundances,
resulting from uncertainties in the stellar parameters and oscillator strengths, are 0.1–0.2 dex.
Possible systematic errors are judged to be of the same order of magnitude. The abundances of
C, N and Ca have been derived from spectrum synthesis, using the C2 band at l ¼ 516.5 nm, the
CN band at l ¼ 388.3 nm (assuming a C abundance as listed above), and the Ca II H þ K lines,
respectively. We measure a carbon isotopic ratio of 12C/13C . 30 from CH A–X lines.
Publishing Group
NATURE | VOL 419 | 31 OCTOBER 2002 | www.nature.com/nature
FeStar=5 10-6 FeO(XFe,O=1.34 10-3)
Æ XFe ~ 6.7 10-9
Æ HE0107-5240 is the most Fe
deficient star presently known
[C/Fe]=+4.02
CStar=5 10-2 CO(XC,§=3.22 10-3)
Æ XC ~ 1.65 10-4
Æ HE0107-5240 is a
typical Globular Cluster star
([M/H]=-1.3)
Standard Scenario
HE0107-5240 was born in a cloud polluted by the ejecta of a
single “standard” supernova
TWO MAJOR INCONSISTENCIES
1.
2.
The Mass Cut required to fit [Ca/Fe], [Ti/Fe] and [Ni/Fe] is
much more internal than the one needed to fit [C/Fe]
The observed ratios among the light elements only, [C/Mg],
[N/Mg] and [Na/Mg], is incompatible with any Mass Cut internal
to the CO core
Int. Structure
[X/Fe]
{X/Mg}
The chemical composition of HE 0107-5240 is INCOMPATIBLE with the
ejecta of a single POPIII Core Collapse Supernova
TWO POPIII Core Collapse Supernovae Scenario
Bonifacio, Limongi & Chieffi 2003 Nature 422,434
Limongi, Chieffi & Bonifacio 2003, ApJL, 594, L126
The most natural solution is to get the Iron peak nuclei from one SN
and the lighter elements from another SN
1.
The observed ratios among just the light elements [C/Mg], [N/Mg]
and [Na/Mg] are perfectly fitted by a 35MO whose mass cut is
located in the He convective shell at M ~ 10 MO
{X/Mg}
2. The star providing the Iron peak nuclei must be less massive because:
i. The yields of the light elements must be negligible
ii. It must have a quite internal mass cut
A 15 MO is the ideal candidate!
Int. Structure
TWO POPIII Core Collapse Supernovae Scenario
Updated Observations
Currently there are 12 stars known to have [Fe/H]<-3.5
The lowest metallicity star known that exhibits an s-process
signature has [Fe/H]=-3.1
Similar abundance pattern of heavy elements
Similar pattern of Light Elements (7 of them)
Large enhancement and scatter of Light Elements (5 of them)
“NORMAL” EMPS
C-RICH EMPS
The 5 stars of the sample sharing a
similar pattern of all the
elements can be represented by
an “Average” (AVG) star
5 of the 12 are C-rich! Is this not unusual
at the lowest metallicities?
Large C, N and O overabundances Æ
heavy element deficient (or Iron
poor) stars rather than metal poor
stars
Is there any way of interpreting the chemical composition
of these stars in terms of ejecta of a population of POPIII
CC-SNe?
Ejecta of single PopIII Core Collapse SNe
M > 25 MO Æ black holes
(heavy element synthesis negligible)
M>25 M
M≤25 M
The chemical composition coming out from the pollution of a population of PopIII
CC-SNe will depend on the combination between the distribution of the various
masses and/or the efficiency of mixing the ejecta of each SN
Ejecta of a population of PopIII Core Collapse SNe
High Mass SNe
Different mixing
efficiency
Low Mass SNe
Gas mainly enriched by the ejecta of High Mass Stars Æ High [Light/Heavy]
Gas mainly enriched by the ejecta Low Mass Stars Æ “Normal” [Light/Heavy]
A different mixing efficiency could explain all the EMPS?
Ejecta of a population of PopIII Core Collapse SNe
M > 25 MO Æ black holes
(heavy element synthesis negligible)
C-rich EMPS
Normal EMPS
Gas mainly enriched by the ejecta Low Mass Stars Æ EMPS ?
Gas mainly enriched by the ejecta of High Mass Stars Æ C-EMPS ?
CONCLUSIONS. II
The element abundance pattern of both the NORMAL and the C-RICH
EMPS can be explained in terms of enrichment of STANDARD POPIII
CC-SNe of different mass and/or different mixing efficiency
Clouds dominated by the ejecta of low mass SNe
Low mass supernovae
produce the same pattern
independent on the mass
Different distributions
and/or mixing degree do
not alter the pattern of all
the elements
NORMAL EMPS
Similar pattern of all the
elements shown by all
Normal EMPS
Clouds dominated by the ejecta of high mass SNe
High mass supernovae
produce a different
pattern of the light
elements depending on
the mass.
Æ
Æ C-RICH EMPS
Different enhancement and
scatter of light elements
shown by C-Rich EMPS
Mixing-Fallback Model (Umeda & Nomoto 2002, ApJ)
Fallback
Mixing region
O
Si
Fe
Explosion
BH
Few days later
Ejected 56Ni mass can be smaller without
changing the Zn/Fe ratio
Mixing-fallback model
Mixing
Fallback
f: ejection factor
Without mixing-fallback model
With mixing-fallback model
MAIN UNCERTAINTIES AND OPEN PROBLEMS
In general the observed abundance patterns of light elements are rather
well fitted by the models (with some discrepancy)
An improvement in the fit between the models and the observations
should be obtained by the computation of a more refined grid of models
On the contrary
The abundance pattern of ALL the iron peak elements are never
reproduced by the models
Cr, Mn and Ni are always very well fitted by the models (Cr and Mn are
made in the same zone by explosive incomplete Si burning)
Sc, Ti, Co and Zn are always heavily underestimated by the models. All
these elements are produced by explosive complete Si burning!
An increase of the C abundance left by central He burning (lower
12C(α,γ)16O cross section) would improve the fit to Sc, Co and Zn
More efficient C burning shell Æ less compact structure Æ higher α rich freezout
Æ increase of [Sc,Co,Zn/Fe]
32S
31P
28Si 29Si 30Si
27Al
24Mg25Mg26Mg
23Na
20Ne 21Ne 22Ne
19F
16O 17O 18O
14N 15N
12C 13C
33S
35Cl
37Cl
34S
36S
32S
31P
28Si 29Si 30Si
27Al
24Mg25Mg26Mg
23Na
20Ne 21Ne 22Ne
19F
16O 17O 18O
14N 15N
12C 13C
33S
35Cl
37Cl
34S
36S
Sic
5 Sii
4
Ox
Nex
3.3
Cx
2.1 1.9
Ye
Mn
Fe
15 M§
Ti
Co
Ti
Ni
Cr
V
Sc
Mass Fraction
Sic
Sii
Ox
Nex
Cx
K
S
Ar
Remnant Ejecta
Si
Ca
Sic
Sii
Ox
Nex
Cx
Cl
He
O
Mg
P
Na
Al
Mass Cut
Chemical
Composition After
The Explosion
Interior Mass (M§)
Ne
C
Element
Production
Site
Sc (45Sc,45Ca)
Co (59Ni)
Ni (58Ni)
Si-cx
Ti (48Cr)
Fe (56Ni)
Si-ix + Si-cx
Cr (52Fe)
V (51Cr)
Mn (55Co)
Si-ix
Si (28Si)
S (32S)
Ar (36Ar)
Ca (40Ca)
Ox + Si-ix
K (39K)
Ox
Ne (20Ne)
Na (23Na)
Mg (24Mg)
Al (27Al)
P (31P)
Cl (35Cl, 37Ar)
C-shell
(hydrostatic
evolution) +
Cx/Nex
He (4He)
C (12C)
N (14N)
O (16O)
F (19F)
Hydrostatic
Evolution
1st Inconsistency
Red Dots = Obtained by choosing the Mass Cut to fit [C/Fe]
Blue Dots = Obtained by choosing the Mass Cut to fit [Ca/Fe]
2nd Inconsistency
2.
The observed ratios among the light elements only, [C/Mg], [N/Mg] and
[Na/Mg], is incompatible with any Mass Cut internal to the CO core
{X/Fe}star=Log(X/Fe)model - Log(X/Fe)star
{C/Mg} - {N/Mg} - {Na/Mg} - solid He – dotted C – dashed Mg