Download Type II supernovae (Inma Dominguez)

Core Collapse SNe
Inma Domínguez
Marco Limongi
 Evolution of Massive Stars
 Hydrostatic Nucleosynthesis
 Explosion Mechanism
 Explosive Nucleosynthesis
 Contribution to the Chemical Evolution
Log Mass Fraction
INTERPRETATION OF THE SOLAR SYSTEM
ABUNDANCES
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
Type II SNe  Chemical Evolution of the Galaxy
Type II SNe  16 < A < 50 and 60 < A < 90
16O
49Ti
60Ni
90Zr
Evolutionary Properties of Massive
Stars:
Progenitors of CCSNe
 M > 12 M
CCSNe
 Central Conditions (T,)
Ignition of ALL
Exothermic
Nuclear Reactions
 The stars is never in degenerate
conditions along its evolution
STELLAR EVOLUTION EQUATIONS
P
Gm

m
4 r 4
r
1

m 4 r 2  ( P, T , Yi )
L
  nuc ( P, T , Yi )  n ( P, T , Yi )   grav ( P, T , Yi )
m
T
GmT

( P, T , Yi )
2
m
4 r P
Mixing-length theory
1 Dimension
Lagrangian
Hydrostatic
STELLAR EVOLUTION EQUATIONS
+ Chemical Evolution
Production + Destruction
Yi
  ci ( j ) jY j   ci ( j, k ) N A  v  j , k Y jYk
t
j
j,k
  ci ( j, k , l )  2 N A  v  j ,k ,l Y jYkYl
2
j ,k ,l
Xi
Ni
Yi 

Ai N A
i  1,........,N
For each time step
1000 (zones) systems of 4+N(isotopes) equations
High Computational Time
HYDROGEN BURNING - PP
4H  He
Proton-Proton Chain
1H
+ 1H  2H + e+ + n
2H
+ 1H  3He + g
3He
PPI
3He
+ 3He  4He + 2 1H
+ 4He  7Be + g
PPII
7Be
7Li
+ e-  7Li + n
+ 1H  2 4He
PPIII
7Be
8B
+ 1H  8B + g
 8Be + e+ + n
8Be
 2 4He
Depending on T the different branchings become active.
In all cases the result is
4 1H  1 4He
HYDROGEN BURNING
CNO Cycle
When C and/or N
12C
and/or
are
present  CNO
+ 1H  13N + g
13N
 13C + e+ + n
13C
+ 1H  14N + g
14N
+ 1H  15O + g
15O
15N
O
CN
 15N + e+ + n
+ 1H  12C + 4He (99%)
16O
+ g (1%)
T  3 107 K
16O
+ 1H  17F + g
17F
17O
 17O + e+ + n
NO
+ 1H  14N + 4He
During the conversion of H into He through the CNO cycle
C and O are burnt and N is produced
Products of CNO
C N O
HYDROGEN BURNING – ENERGY GENERATION
The CNO cycle is more efficient than he PP chain over a certain Tcritica
CNO
PP
From Hydrostatic Equilibrium Eq:
M
P  T  3 T
R
Tc 
M
R
Massive stars H-burning
M2
Pc  4
R
P
Gm

m
4r 4
Central Temperatura
scales with Total Mass
CNO
cycle
HYDROGEN BURNING - CONVECTIVE CORE
The Energy generated by the CNO-cycle depends strongly on T
High Energy Flux  Increases Radiative Gradient  A Convective core Develops
3
 dT 
F
  
3
4acT
 dr rad
Masssive stars burn H within a Convective core
At high T the main contribution to the Opacity comes from the Thomson Scattering
Th  0.2  (1  X H )
When the H decreases, the Opacity decreases and the Convective Core receeds
and finally, at H-exhaustion, disappears
HYDROGEN BURNING
–
Ne-Na, Mg-Al Cycles
If during the central convective H-burning T are high enough log T=7.5-7.8
 Active Ne-Na e Mg-Al cycles
Ne-Na Cycle
20Ne
+ 1H  21Na + g
21Na
21Ne
Mg-Al Cycle
24Mg
 21Ne + e+ + n
25Al
+ 1H  22Na + g
22Na
+ 1H  25Al + g
25Mg
 22Ne + e+ + n
+ 1H  26Al + g
26Al
22Ne
+ 1H  23Na + g
26Mg
23Na
+ 1H  20Ne + 4He
27Al
 25Mg + e+ + n
 26Mg + e+ + n
+ 1H  27Al + g
+ 1H  24Mg + 4He
Final results of the operation of these cycles Na-Na e Mg-Al

21Na
&
25Mg

22Ne
is reduced by a factor 2

23Na
&

26Al

20Ne, 24Mg
26Mg
practically burnt
increase by a factor 6 & 2, respectively
produced (~10-7)
&
27Al
do not change
STRUCTURE AT CENTRAL H-EXHAUSTION
The He-core is much more dense than the H-envelope because the mean
molecular weight for 4He is greater than for 1H
 Matter within the He-core is more compact

1
Zi  1
 A
i
i
 H  0.5  He  1.333 C  1.714 O  1.778
He core
H envelope
The synthesis of heavier isotopes increases the mean molecular
weight and the structure becomes more compact
HYDROGEN SHELL BURNING
 At central H-exhaustion  H-burning sets in a Shell outside the He-core.
 HR diagram: the star moves to the red
 A convective envelope forms, the inner border of this envelope reachs zones
chemically modified by he central H-burning.
 The 1st dredge-up occurs: material processed by nuclear reactions is
transported to the surface
H exhaustion
Convective envelope
Start Conv. Env.
H burn.shell
dup
He core
He
H conv. core
He conv.
core
He burn. CO core
shell
HELIUM BURNING – 3
At central H-exhaustion, the He core is mainly composed by
4He
(98%) & 14N (1%)
Withouth Nuclear Energy generation within the core, it contracts and
Tc increases
When Tc ~ 1.5 108 K  Efficient He-burning
At the beginning 4He  8Be and 8Be rapidly decays to 4He
4He
+ 4He  8Be + g
8Be
 4He + 4He
Later, at higher T and   the equilibrium abundance of 8Be increases
and so increases the probability of the reaction
8Be + 4He producing 12C
4He
+ 4He  8Be + g
8Be
8Be
 4He + 4He
+ 4He  12C + g
3 4He  12C + g
HELIUM BURNING – REACTIONS
Initially: 4He in 12C
But when 12C abundance is significant and 4He abundance is reduced,
it is more likely that 4He is captured by 12C than by 4He:
3 4He  12C + g
12C
+ 4He  12O + g
16O
+ 4He  20Ne + g
20Ne
The first 2 reactions
are more efficient
+ 4He  24Mg + g
3 4He Nuclear Cross Section depends markedly on T
Like H-burning (CNO cycle)
He-burning occurs within a convective core
HELIUM BURNING: s-process
14N
14N
produced by the CNO cycle
+ 4He  18F + g
18F
18O
+
22Ne

4He
18O

+
e+
22Ne
+n
78Rb
86Kr
+g
77Kr
+ 4He  25Mg + n
85Br
76Br
79Rb
80Rb
81Rb
82Rb
83Rb
84Rb
85Rb
80Kr
81Kr
82Kr
83Kr
84Kr
80Br
81Br
82Br
83Br
78Se
79Se
80Se
81Se
82Se
77As
78As
79As
80As
81As
75Ge
76Ge
77Ge
78Ge
79Ge
80Ge
74Ga
75Ga
76Ga
77Ga
78Ga
79Ga
87Kr
88Kr
78Kr
79Kr
86Br
b-
87Br
77Br
78Br
79Br
b84Se
75Se
83As
74As
85Se
86Se
76Se
77Se
84As
b-
85As
75As
76As
b
73Ge
74Ge
72Ga
73Ga
n,g
In Massive  during central He-burning, elements
heavier than Fe are synthesized by the s-process.
s-process depends on free neutrons and the neutron abundance depends on Z 
The final s-element abundances scale with initial metallicity
HELIUM EXHAUSTION
The most abundant isotopes at
central He-exhaustion:
12C
16O
20Ne
25Mg
26Mg
12C
The first three are produced by:
22Ne
25Mg
3 4He  12C + g
12C
26Mg
+ 4He  12O + g
ex He c.c.
25Mg
16O
+ 4He  20Ne + g
&
26Mg
14N
come from the
20Ne
H sh.
Conv. Envelope.
Core di
CO
+ 4He  18F + g
18F
18O
14N-chain
16O
 18O + e+ + n
+ 4He  22Ne + g
4He

22Ne
+
22Ne
+ 4He  26Mg + g
12C/16O
25Mg
+n
12C, 16O, 20Ne, 25Mg
& 26Mg are the most
abundant isotopes and are produced by He-burning
with the surface abundance
ratio depends on the 12C + 4He  12O + g nuclear cross section
that it is still NOT well known at the energies of the He burning.
This ratio has a strong influence on the subsequent evolution
HELIUM EXHAUSTION: s-process elements
The most abundant
elements are:
70Ge
80Kr
70Ge
74Se
and
80Kr
74Se
ex He c.c.
Core di CO
Heavier nuclei, like
87Rb, 88Sr, 89Y, 90Zr
are not expected to
be produced
H sh.
Conv. Envelope.
HELIUM SHELL BURNING – CONVECTIVE SHELL
At central He exhaustion, He burning moves to a shell just outside the CO core
The following evolution is characterized by the development of a convective He-burning shell
limited by the CO core and by the H-burning shell.
The chemical composition of this shell, that will be active till the collapse, tends to get
frozen because the evolution of the star is more and more rapid at the advanced phases.
Convective envelope
H burn.shell
dup
He conv.shell
He core
He
H conv. core
He conv. core He burn.
shell
CO core
STRUCTURE at He-exhaustion
At central H-exhaustion, the  is composed by a CO core, a He-shell
and a rich H envelope
CO core
He core
H envelope
The two density gradients correspond to the border of the He core
(~ 9 M) and to the border of the CO core (~ 6 M )
This density profile is important for the explosion properties
ADVANCED EVOLUTIONARY PHASES: NEUTRINO
DOMINATED
Now the CO core, produced by the central He-burning, contracts
During the contraction the  and T within the core favours the production
of thermal neutrinos produced by pair anhilation.
At T>109 K high energy photons produce e+e- pairs
g  e  e
That suddenly recombine to produce a photon.
BUT once over 1019 times, e+e- produces a neutrino-antineutrino pair
g  e   e   (1 / 1019 times)  n e n e
This energy sink increases along the subsequent phases
up to the pre-collapse phase
Advanced evolutionary phases of massive stars are called
“neutrino dominated”
ADVANCED EVOLUTIONARY PHASES:
NEUTRINO LUMINOSITY
From now on the energy
losses:
Photons from the surface
Neutrinos from the center
108
Nuclear
Neutrino
Photon
Up to C central ignition the main energy losses are due to photons and
after are due to neutrinos. As the nuclear energy gives the star what is
lossing, it follows first the luminosity of photons, and after, the neutrino
luminosity
EVOLUTIONARY TIMES
Enuc is the energy per gram coming from nuclear reactions,
If this is the only energy source in a star of mass M:
L
Enuc
M
tnuc
Nuclear time scale:
M
tnuc  Enuc
L
H burning: 4 1H  4He
DM = 4 x 1.0078 – 4.0026 = 0.0287 AMU = 0.0287/4 AMU/nucleon = 0.007 AMU/nucleon
Enuc = 0.007 x 931.1 x 1.602
1 AMU = 931.1 MeV :
He burning: 4 4He 
10-6 x 6.022 1023 = 6.44 1018 erg/g
1 MeV= 1.602 10-6 erg
:
NA = 6.022 1023 nucleon/g
16O
DM = 4 x 4.0026 – 15.9949= 0.0115 AMU = 0.0115/16 AMU/nucleon = 0.0009 AMU/nucleon
Enuc = 0.0009 x 931.1 x 1.602 10-6 x 6.022 1023 = 8.70 1017 erg/g
O burning: 2
16O

32S
DM = 2 x 15.9949 – 31.9720= 0.0177 AMU = 0.0177/32 AMU/nucleon = 0.0005 AMU/nucleon
Enuc = 0.0005 x 931.1 x 1.602 10-6 x 6.022 1023 = 4.98 1017 erg/g
For fix mass, Luminosity and amount of fuel
t He
tO
 0.13
 0.08
tH
tH
From models:
tHe
t
 0.11 O  5.56  108
tH
tH
The luminosity increases drastically due to neutrino losses 
The evolutionary times are drastically reduced
!
Advanced burning stages
Neutrino losses play a dominant role in the evolution of a massive star
beyond core He burning
At high temperature (T>109 K) neutrino
emission from pair production start to
become very efficient
g
g
n
g
n
n
n
n
n
g
g
n
g  e  e  n e  n e
g
g
n
g
t  Enuc
M
L
Evolutionary times
reduce dramatically
CARBON BURNING
Central C combustion stars ~104 years after central He-exhaustion
Tc ~ 7 108 K e c ~ 1 105 g/cm3
C-burning depends on the
12C/16O
ratio left after central He burning,
12C(,g)16O
on the amount of fuel
The formation of a Convective Core depends on the existence of a
positive energy flux
12C
abundances
determines the nuclear
energy generation rate
nuc > n
A Convective Core
develops
nuc < n
NO Convective
Core
In general, for a fix 12C(,g)16O reaction rate and mixing technics
12C abundance decreases for higher initial masses
In the 25M  central carbon combustion occurs in radiative conditions
Synthesis of Heavy Elements
At high temperatures a larger number of nuclear reactions are activated
Heavy nuclei start to be produced
C-burning
T ~ 109 K
Ne-burning
T ~ 1.3 109 K
Synthesis of Heavy Elements
O-burning
T ~ 2 109 K
Synthesis of Heavy Elements
At Oxygen exhaustion
Balance between forward and
reverse reactions for increasing
number of processes
T ~ 2.5 109 K
c + d
a + b
At Oxygen exhaustion
At Si ignition
At Si ignition
(panel a + panel b)
T ~ 2.5 109 K
T ~ 3.5 109 K
T ~ 3.5 109 K
A=44
Sc
Si
Equilibrium
A=45
56Fe
28Si
Equilibrium
Partial Eq.
Out of Equilibrium
Eq. Clusters
Out of Eq.
56,57,58Fe, 52,53,54Cr,
55Mn, 59Co, 62Ni
MATTER PROPERTIES AT HIGH TEMPERATURE :NSE
The chemical composition of matter in NSE is a function of T  Ye
Yi NSE  f (T ,  ,Ye )
When the neutronization changes
  1 2Ye
Ye  
i
Zi
Xi
Ai
The nuclei with that neutron excess are favoured (with higher binding energies)
T  5 10 K   10 g/cm
9
8
3
=0.000,Ye=0.5000, 56Ni
=0.038,Ye=0.481, 54Fe
=0.072,Ye=0.464, 56Fe
=0.104,Ye=0.448, 58Fe
58Ni, 53Mn
O Conv. Shell
28Si, 32S, 36Ar,
40Ca, 34S, 38Ar
C Conv. Shell
20Ne, 23Na,
24Mg,25Mg, 27Al
s-process
+
16O
H Centrale
H Shell
He Shell
4He
1H
28Si
“Fe”
He Centrale
16O, 12C
20Ne
He Shell
16O, 12C
12C
H Centrale+Shell
14N, 13C, 17O
+ sprocess
He Centrale
54Fe, 56Fe, 55Fe,
C conv. Shell
Si Burning
Main Products
O conv. Shell
Burning Site
Si burning(Cent.+Sehll)
PRE-SUPERNOVA MODEL: CHEMICAL COMPOSITION
Studying the different isotope abundances in detail is possible to know from which
burning phase they come from or the interior region of the star where they were produced
PRE-SUPERNOVA MODEL: Fe-CORE STRUCTURE
Fe/Si
Si/O
CO/He
He/H
16O
“Fe”
28Si
20Ne
12C
EXPLOSION
The gravitational collapse of a stars with M  12 M could liberate an energy of
Ebind
3
1
2 1
 GM     1053 erg
5
 R2 R1 
R2  10Km
Most of this energy increases the electron energy and, after electron captures,
is converted in neutrino energy
Just a small fraction is used to eject (kinetic energy) the envelope
So, the key question is to find a mechanism able to transform a
small fraction of the binding energy left during the collapse in kinetic
energy of the envelope with the observed velocities ( 104 km/s)
Explosive Nucleosynthesis and Chemical Yields
Hydrocode (Collella & Woodward 1984)
• Explosive Nucleosynthesis: same
nuclear network adopted in the
hydrostatic evolutions
28Si
“Fe”
H Central
H Shell
He Shell
He Central
C conv. Shell
Piston
• Explosion: 1D PPM Lagrangian
Si burning
The explosion can be simulated by means of
a piston of initial velocity v0, located near
the edge of the iron core
O conv. Shell
Explosion Mechanism Still Uncertain
4He
16O
20Ne
12C
v0 is tuned in order to have a given amount of 56Ni ejected and/or a
corresponding final kinetic energy Ekin
1H
EXPLOSIVE NUCLEOSYNTHESIS
Passing through the envelope the Shock Wave increases the density and
temperature and nuclear reactions occur
We may define the burning time-scales for the available fuels :
Si, O, Ne, C, He and H
Y
i  
Y
These time scales are determined by the corresponding destructive reactions
 i  f (T ,  )
Assuming the explosion time ~1s
T  4  109 K Si burning
T  3.3  109 K O burning
T  2.1109 K Ne burning
Burning products are similar to those
obtained in hydrostatic burning
T  1.9 109 K C burning
  105 g/cm3
He burning
He-explosive burning is not efficient in SNII
EXPLOSIVE NUCLEOSYNTHESIS
Analyzing the most eficient processes:
Still out of NSE: Products are similar to those from hydrostatic
burning
T  3.3  109 K
T  1.9  109 K
EXPLOSIVE CARBON BURNING:
Products:
20Ne, 23Na, 24Mg,25Mg, 26Mg
T  2.1  109 K
EXPLOSIVE NEON BURNING:
Products:
T  3.3  109 K
16O, 24Mg
Starting NSE
+
27Al, 29Si, 30Si, 31P, 35Cl, 37Cl
(direct and inverse process)
EXPLOSIVE OXYGEN BURNING:
56Fe
A=45
A=44
28Si
T  4.0  109 K
2 clusters at quasi-NSE separated by
A44.
No connection between the 2 clusters
Clusters di equilibrio
YiQSE  f (T ,  ,Ye , Yi )
Q
Products:
28Si, 32S, 36Ar, 40Ca
+
34S, 38Ar
EXPLOSIVE NUCLEOSYNTHESIS
EXPLOSIVE INCOMPLETE SILICON BURNING:
56Fe
A=45
T  5.0  109 K
At this T the 2 clusters connect at A44.
Most of the matter A<44  just part of
the upper cluster
A=44
28Si
28Si
reachs
Yi  f (T ,  ,Ye ,28Si)
Products:
Clusters di equilibrio
36Ar, 40Ca
+ 56Ni(56Fe), 54Fe,
52Fe(52Cr),51Cr(51V), 55Co(55Mn), 57Ni(57Fe), 58Ni
EXPLOSIVE COMPLETE SILICON BURNING:
T  5.0  109 K
At this high temperature: NSE !!!!!!
All 28Si is burnt to Fe-peak elements.
Abundances depend on neutronization !!
For NZ 56Ni is the most abundant nuclei
Yi  f (T ,  ,Ye )
Products: Iron Peak Nuclei
Full NSE
EXPLOSIVE NUCLEOSYNTHESIS
Changes in T and  following expansion are crucial for the nucleosynthesis
During the explosion Temperatures are very high
It could be assumed that matter behind the shock is radiation dominated
Eexpl
4 3 4
  R aT
3
R, T= Location and T of the shock
The shock propagates in all directions (sphere)
Each radial coordinate in the presupernova
model will reach a maximum temperature
 3Eexpl 

Tmax  
3

4

R
a
PSN 

1
4
EXPLOSIVE NUCLEOSYNTHESIS
Complete
Si burning
Incomplete Explosive
Si burning
Oxygen
5.0  109 K
4.0  109 K
3.3  109 K
NSE
QSE
1cluster
QSE
2cluster
Sc,Ti,Fe,
Co,Ni
Cr,V,Mn,
Fe
3700
Si,S,Ar,
K,Ca
5000
Explosive
Neon
Explosive
Carbon
2.1  109 K
1.9  109 K
Mg,Al, P, Cl
6400
Ne,Na,Mg
11750
Untouched Zone
For Eexpl=1051 erg we could infer in the presupernova model which regions (volumes)
experience each burning
13400
EXPLOSIVE NUCLEOSYNTHESIS: PROGENITOR
Influence of the Progenitor:
1) M-R RELATION (= density profile):
Fix the mass inside a certain volume
2) Ye (neutronization):
In those zones that reach NSE or QSE determines the rate between
protons and neutrons
T=5 109 K, = 108 g/cm3, Ye=0.50  56Ni=0.63 – 55Co=0.11 – 52Fe=0.07 – 57Ni=0.06 – 54Fe=0.05
T=5 109 K, = 108 g/cm3, Ye=0.49  54Fe=0.28 – 56Ni=0.24 – 55Co=0.16 – 58Ni=0.11 – 57Ni=0.08
3) Chemical Composition :
For those zones that experience normal burnings (ie. Explosive
Carbon e Neon burnings) fix the amount of fuel available.
MASS CUT
During the explosion internal zones fall back.
At some point part of the matter is Expanding
and some Collapsing
Depending on v compare to vesc 
The mass coordinate at the bifurcation is
defined as the Mass Cut
The Mass Cut depends on the piston initial
velocity
 M pist , v0  Ekin , M cut
In general, for greater initial velocities
Smaller Mass Cut
Greater kinetic Energies
1.110
1.144
1.170
1.220
1.250
Mass Cut
1.263
The lack of a explosion model makes the MASS CUT and the KINETIC ENERGY quantities that
depend on parameters (initial energy or piston initial velocity and place at which the explosion is
started)
EXPLOSION PROPERTIES: CHANGES IN CHEMESTRY
Taken:
v0=1.5550 109 cm/s
Mcut=1.89 M
Ekin=1.14 foe
Mass Cut
Pre = Dotted
Post = Solid
Si-c
4He
Si-i Ox
Nex
Cx Untouched
16O
28Si
20Ne
1H
12C
 The changes in composition due to the explosion occur only at the most internal ~3.1 M
 Outside the chemical composition remains untouched. It is that from the hydrostatic
burning
 The complete explosive Si burning and part ot the incomplete explosive Si burning fall
back to the compact remant
MASS CUT CALIBRATION: LIGHT CURVES
From the LC we obtain information for the Mcut
After an initial phase, different for the different types of SNe, the LC is powered by the photons
produced by the radioactive decay
56
8.8
111
Ni  56Co  56Fe
Total
56Ni
56Co
Based on the Bolometric LCs and on the
distance, we can deduce the amount of
56Ni produced during the explosion
56Ni=0.15
M
56Ni=0.07
M
56Ni=0.01
M
56Ni
is produced in the most internal zone
depends critically on the Mass Cut

The Mass Cut may be choose to reproduce a
certain amount of 56Ni in agreement with the
observations.
The theoretical kinetic energy must be compatible
with the observed
MASS CUT CALIBRATION vs INITIAL MASS
From the observed initial mass of the progenitor we may obtain an empirical relaction between
this mass and the 56Ni produced (or Mcut)
M cut  f ( M i )
M (56Ni )  f ( M i )
Hamuy et al. 2003
PROBLEMS !!!!
 Few estimations of the progenitor initial mass from the observations
 Similar masses give very different
56Ni
masses
CHOOSING A MASS CUT
1) FLAT Case: All masses produce the same 56Ni mass = 0.05 M  For each
model a different mass cut is chosen in order to reproduce this amount of Ni
2) TREND Case: We adopt a relation between Initial Mass and
Mi
(M)
M(56Ni)
(M)
13
0.15
15
0.10
20
0.08
25
0.07
30
0.05
35
0.05
56Ni
Mass:
PRODUCTION FACTORS
To compare with Solar Abundances we introduce the Production Factor
Mto t
PFi 
 X i dm
Mcu t
Mto t
S un
X
 i dm
Mcu t
Two isotopes with the
same Production Factor
Same Rate as in the Sun
Oxygen is produced only by Type II SNe and is the most abundant
element produced by SNII  Oxygen Production Factor is a Good
Metallicity indicator
It is useful to normalize all PF to that of Oxygen to
show wich isotopes follow Oxygen (Z)
INTEGRATED YIELDS (Elements)
Yields from 13-35 M + Salpeter Mass Function
It is assumed that all masses produce the same amount of
56Ni
(FLAT)
We consider “Solar Scaled” with respect to O all elements with a PF within a factor 2 of the O PF
Dots: 13 – 15 – 20 – 25 – 30 – 35 M
Solid line: Salpeter Mass Function
Flat
56Ni
=> 0.05 M
The yields produced by a generation of massive stars integrated by a Salpeter IMF
depend mainly on the yields coming from a 20-25 M star
Contribution of Type Ia SNe
Production of Fe  the percentage of SNIa, relative to SNII, has been
fixed by requiring that PFFe=PFO
Open circles = No SNIa Filled circles = 12% SNIa
1) SNIa contribute only to the
Solar System abundances of
nuclei in the range Ti-Ni
2) The inclusion of SNIa brings
50Ti and 54Cr into the band of
compatibility  50Ti and 54Cr
become scaled solar compared
to O
3)
14N
and lot of
heavy elements come
from AGB stars
CONCLUSIONS
with mass loss: 11 -120 M
 Massive Stars are responsible for producing elements from
12C (Z=6) up to 90Zr (Z=40)
+ r-elements
 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
Pre/Post SN models and explosive yields available at
http://www.mporzio.astro.it/~limongi
Alessandro Chieffi & Marco Limongi (ApJ 1998-2007)
Uncertainties in the computation PreSN Models
 Extension of the Convective Core (Overshooting, Semiconvection)
 Mass Loss
Uncertainties in the computation of the Explosion Models
 Explosion itself
Piston:
 Mass-cut - Mini

56Ni
(LC)
 Energy (vexp)
Navegamos sin rumbo
a través del obscuro Océano Cósmico
Estrellas y planetas
en un espacio infinito…
¿ Tiene sentido nuestra
presencia en el Universo ?
¿ Podemos ganar
la liga de campeones ?
IDEAL
ORCEMAN
by C. Hernández
CHEMICAL ENRICHMENT BY A GENERATION OF MASSIVE STARS
The 25 M solar model could be considered as the “typical” case, representative of
stars from 13 to 35 M
If we compute the YIELDS (ejected abundances in solar masses) of the different
isotopes produced by a grid of models (~13 to 35 M), we could compute the
chemical contribution of a generation of Massive Stars to the ISM
These YIEDS are ingredients in a Chemical Evolution Model for the Galaxy, includes SFR,
IMF & Infall
In principle, the chemical solar distribution is a consequence of different generations
of stars with different initial compositions
The metallicity of the ISM is expected to increse continously and with longer timescales than the evolutionary time of the stars that contributes to the chemical
enrichment
We expect that the YIELDS of a generation of masive solar metallicity stars
explain the solar distribution
Integrated Yields adopting a different Mi-M(56Ni)
relation
13
Flat
56Ni
=> 0.05 M
Trend
56Ni
15
20
25
30
35 M
=> 0.15–0.10–0.075-0.07–0.05–0.05 M
Int. Mass Stars
n processes
Int. Mass Stars
The only elements that vary between case “Flat” and case “Trend” are Fe and Ni and,
at a smaller extent also Ti, Co and Zn (i.e. elements produced in the deep layers of the
exploding mantle)
The majority of the elements have PFs compatible with that
of O  show a scaled solar distribution
The Final Fate of a Massive Star with mass loss: 11 -120 M
Z=Z
E=1051 erg
SNII
SNIb/c
WNL
WNE
WC/WO
Fallback
Mass (M)
RSG
Black Hole
Neutron Star
Initial Mass (M)
Limongi & Chieffi, 2007
Individual Yields
Different chemical composition of the ejecta for different masses
Averaged Yields
Yields averaged over a Salpeter IMF
 (m)  m
  2.35
Global Properties:
Initial Composition
(Mass Fraction)
X=0.695
Y=0.285
Z=0.020
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)
Observed MPro smaller than LC models
predict
Li et al.
Smartt et al.
van Dyk et al.
Initial Mass Function
mu~ 100 M; ml ~ 0.1 M
mrem  Stellar evolution
IMF ≈ Present Day MF
for massive stars
IMF ...universal?
Definitions
AMU (atomic mass unit, mu) 
1/12 mass of
12C
muc2 = 931.478 MeV
Cross section: Probability per pair of particles of occurrences of a reaction
  cm2
n  cm3 /s