Transcript Introduction - Departamento de Astronomía

The Physics of
Supernovae
Inma Domínguez
Universidad de Granada
Santiago de Chile, octubre de 2007
SN 1987A
 Chemical Evolution
 Cosmology
 Trigger Star formation
 Neutrinos
 BH, NS, GRBs
 Reionization of the Universe
etc etc
Supernovae are one of the most
energetic explosive events in Nature
• BRIGHT
A SN in 10 sec releases 100 times the energy that
the sun releases in all its life
SN1054 was as luminous as the moon for some days
• RARE: About 1 per century in our Galaxy
Last recorded seen by naked-eye :1006
(Lupus), 1054 (Chinese), 1572 (Brahe),
1604(Kepler)
• BRIEF: Luminosity falls by a factor of 100 in
4 months
Standard Candles
Fainter  Further
 Distance Modulus
 Luminosity Distance
SNe Classification
Based on spectra and light curve morphology
II P
Type II
II L
SNe
I b (strong He)
I c (weak He)
Core
collapse
of
massive
stars
Type I
I a (strong Si)
Thermonuclear
explosion
Basic SN type spectra
Light Curves
Type Ia SN
•Similar luminosity
•Similar spectral evolution
 Good distance indicators
Cosmological parameters
Type II SN
•Dramatic differences
•II-P (plateau)
•II-L (rapid declination)
• Cosmology
SNe RATE
SN rate per unit Mass (10-10 M 10-2 yr (Ho/75)2
Galaxy Ia
Ib/c
II
E-S0
0.04
< 0.01
< 0.01
S0a-Sb
0.065
0.026
0.12
S0c-Sd
0.17
0.067
0.74
Irr
0.77
0.21
1.7
Mannucci et al. 2005
SN Ib/c & SNII Absent in E-S0
Young populations
Short lived progenitors
Massive 
SN Ia
in E-S0 Old populations
Long lived progenitors
Low mass  in Binary Systems
SN Ia rate  in Spirals Galaxies-with SFR
Part of SN Ia comes from a younger population
Cappellaro et al 2003, Mannucci et al. 2005, Sullivan et al. 2006
Stellar Evolution
M<0.8 M
t>1/Ho
30 Myr<t< 15 Gyr
0.8<M/M<8
0.5<Mf /M<1.1
AGB
SN Ia
CO WD
t~10-30 Myr
8<M/M<11
Mf =1.2-1.3 M
ONeMg WD

t~ 1-10 Myr
11<M/M<100
Mf =1.2-2.5 M
SN II Ib/c
Fe collapse NS/BH
M>100 M
t~
1 Myr
may or may not
explode
Classification of SNe
~ 4000 SNe
(nowadays > 300 /yr)
Solar System Abundances
1
0 H4
He
-1
16O
12C
O
20Ne
Log Mass Fraction
-2
-3
56Fe
-4
-5
N=50
-6
N=126
N=82
-7
-8
-9
-10
-11
-12
0
20
40
60
80
100
120
Atomic Weight
The most abundant isotopes:
 1H
 4He
 16O
12C
20Ne
(-elements)
140
160
180
200
220
50 yrs !!
Origin of the Elements: Inside the Stars
Observational Evidences:
 Pop II 
Less heavy elements by a factor of 100
Our Galaxy has synthesized 99 % of the
heavy elements during ¡ts evolution
 Merril (1952) discovered Tc in 
All Tc isotopes decay t1/2  106 yr
Tc has been synthesized inside the star
Origin of the Elements:
Nuclear Statistical Equilibrium (NSE) ?
 Klein, Beskow & Treffenberg
(1947)
Studied the abundances at NSE
in function of T and 
rate nuc. re. = inverse rate
N ( A, Z )  f (nn , T )
This mechanism could not reproduce the observed
abundances

But NOT bad for the Fe peak !!
Binding Energy per nucleon
BE/c2=[Zmp + (A-Z)mn - m(A,Z)]
BE/A
56Fe
© Rolfs & Rodney 1988
smallest
mass per nucleon

to 56Fe
exothermic
reactions
The interpretation of the abundances
 The Peaks in the abundances of
4He, 12C, 16O, 20Ne
and other  elements
 capture nuclear reactions inside the stars
 Fe-peak elements
56Fe
is the isotope with higher binding energy
56Fe is the last product of exothermic nuclear
fusion reactions, NSE
 Elements heavier than Fe
High Coulomb barrier for charge reactions
Neutron captures
Most
abundant
nuclei
 Nuclear Physics
 Physical Conditions
 Where & When ??
Anders & Grevesse 1989
Solar System Abundances
© Cameron 1982
Abundances peak at the “magic numbers”,Z: 2, 8, 20, 28, 56, 82
He, O, Ca, Fe, Ba, Pb
The familiar picture
 H burning (the most effective, with an average of 7MeV per
nucleon of generated energy): produced 4He, 3He, and gives
(generally secondary) contributions to intermediate nuclei up to Si.
 He burning (the second-most effective): produces 12C,
some 20Ne, plus secondary chains starting from 14N or
leading to neutron generation.
16O,
13C
and
 Fusion of intermediate nuclei - 12C, 16O, 20Ne, 28Si
 nuclei below and up to the Fe-peak.
 Nuclear statistical equilibrium (NSE) processes, crossing the
peak at 56Fe - 56Ni.
 Explosive nucleosynthesis, starting from NSE and reorganizing
abundances up to 65Cu, occur in CCSNe and in SN Ia.
 Neutron captures (slow and rapid – s and r - processes).
Solar System Abundances
BBN
AGB
SNII
SNII
SNII ?
SNIa
BBN
Anders & Grevesse 1989
Cameron 1982
AGB
Some definitions…
• “Metals”: elements heavier than helium, Z
• “Metallicity”: [Fe/H] = log (Fe/H) – log (Fe/H)
• “Abundance ratio”: [X/Y]= log (X/Y) – log (X/Y)
* Abundance scale by number:
* Mass fractions:
X= Hydrogen (X~0.71)
Y= Helium 4 (Y~0.27)
Z= Metals
(Z~0.02)
 Population I
12  log N(H)
X+Y+Z= 1
objects (stars): Z ~ Z
 Population II : Z << Z
 Population III : Z ~ 0 (not detected yet ?)
Stellar Evolution & Nucleosynthesis
 The activation of a nuclear burning phase
 The stellar life-time
DEPEND on
 AGB
Mass
Planetary Nebulae
White Dwarfs
(if) Binary Systems
 Novae
 SNe Ia
 AIC: Neutron 
(Pulsars)
M
Tc 
R
 CCSNe
 Neutron 
(Pulsars)
 Black Holes
“Less” in Z…
Low mass stars M < 8 M
AGB/Planetary Nebulae return
C, N, s-elements etc to the ISM
Exploding CO WDs
(accreting mass from a companion)
Type Ia Supernovae
(SN Ia or
Thermonuclear SNe)
SN Ia produce
~2/3 of the observed
Fe in the Universe
Massive 
25 M
Chieffi, Limongi, Straniero 1998
Massive stars M ≥ 8-10 M
Core Collapse Supernovae eject O, Mg, Ti and
likely r-p-elements into the ISM
Log Mass Fraction
Origin of the elements
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. = ν induced reactions in SNII;
IMS = Intermediate Mass Stars; SNII = Core collapse supernovae;
SNIa = Thermonuclear supernovae; s-r = slow-rapid neutron captures
The Origin of the Elements up to Zn
ApJS 1995
L* M < 8M
 neut. Irra
CR Cosmic Rays
s shell
x Explosive
 rich
freeze out
Yields
 Low and Intermediate Mass Stars
4He
C N s-process (A > 90) elements
Lattanzio et al., Meynet & Maeder, Marigo et al., Siess et al.
Straniero et al. (TERAMO), Siess et al., Van den Hoeck & Groenewegen
Ventura et al.
 Type Ia Supernovae
Fe and Fe-peak
Nomoto et al., Iwamoto et al. Höflich et al., Thielemann et al.
 Massive stars
-elements (O, Ne, Mg, Si, S, Ca),
some Fe-peak, s-process elements (A < 90)
and r-process elements.
Woosley & Weaver / Limongi & Chieffi (ORFEO)
Some definitions
 Yields
Yield i 
Mass Loss !!
t
0
(
X
X
 i i )dm
Meje
in M
 Production Factor
 X dm
i
PFi 
Meje
0
X
 i dm
Meje
Yields + Evolution-Time 
Chemical Evolution
SN II
-elements
SN Ia + SNII
Fe

20Ne
24Mg
28Si
32S
36Ar
time
Mg / Fe  log Mg / Fe - log Mg / Fe
Chemical Evolution
40Ca
-enhancements appear naturally due to the different
life-times between SNII and SNIa… but at what level?
and when?
Modification of the IMF: more massive stars produce more “alphas”
Modification of the SFR: more “alphas” produced before SNIa appear
© McWilliam (1997)
Ingredients of GCE
Initial conditions
Big Bang abundances
Prompt initial enrichment
Initial mass function (IMF)
Relative birthrates of stars with different masses
Star formation rate (SFR)
Constant, burst, interruptions etc
Stellar yields vs. stellar mass and metallicity
SNII, SNIa, AGB, Novae, etc
Galactic gas inflow/outflow
Late infall of primordial gas etc
Supernova-driven galactic winds etc
Stellar & gas dynamics
STELLAR EVOLUTION EQUATIONS
P
Gm
m
4 r 4
r
1

m 4 r 2  ( P, T , Yi )
1 Dimension
Lagrangian
Hydrostatic
L
  nuc ( P, T , Yi )   ( P, T , Yi )   grav ( P, T , Yi )
m
T
GmT
dlnT
( P, T , Yi )

2
dln P
m
4 r P
Yi
  ci ( j ) jY j   ci ( j, k ) N A < v > j , k Y jYk
t
j
j,k

2
c
(
j
,
k
,
l
)

N A < v > j ,k ,l Y jYkYl
i
2
j ,k ,l
i  1,........,N
+ Chemical
Evolution
STELLAR EVOLUTION
EQUATIONS
Convection (a problem !!)
t mix  t
 Time-dependent convection
t mix  t nuc
 Mixing-Nuclear burning coupled
Micro-physics
 EOS




Opacity
Nuclear Cross Sections (Strong & Weak)
Screening factors
Neutrinos
NUCLEAR
NETWORK
60Zn 61Zn 62Zn 63Zn 64Zn 65Zn 66Zn 67Zn 68Zn
High number of Isotopes
High Number of Nuclear Reactions
57Cu 58Cu 59Cu 60Cu 61Cu 62Cu 63Cu 64Cu 65Cu 66Cu 67Cu
56Ni
57Ni
58Ni
59Ni
60Ni
61Ni
62Ni
63Ni
64Ni
65Ni
54Co 55Co 56Co 57Co 58Co 59Co 60Co 61Co 62Co
52Fe
53Fe
54Fe
55Fe
56Fe
57Fe
58Fe
59Fe
60Fe
61Fe
51Mn 52Mn 53Mn 54Mn 55Mn 56Mn 57Mn
p, n and  captures
e± captures
b± Decay
48Cr 49Cr 50Cr 51Cr 52Cr 53Cr 54Cr 55Cr
41Sc
42Sc
45V
46V
47V
48V
49V
50V
51V
52V
44Ti
45Ti
46Ti
47Ti
48Ti
49Ti
50Ti
51Ti
43Sc
44Sc
45Sc
46Sc
47Sc
48Sc
49Sc
50Sc
40Ca 41Ca 42Ca 43Ca 44Ca 45Ca 46Ca 47Ca 48Ca 49Ca
37K
38K
39K
40K
41K
42K
35Ar 36Ar 37Ar 38Ar 39Ar 40Ar 41Ar
1H
3He
4He
2H
3H
25Al
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
(,n)
(p,n)
b-,
23Mg 24Mg 25Mg 26Mg 27Mg
n
21Na 22Na 23Na 24Na
(p,g)
(,p)
20Ne 21Ne 22Ne 23Ne
7Be
8Be
6Li
7Li
17F
18F
19F
20F
15O
16O
17O
18O
19O
13N
14N
15N
16N
12C
13C
14C
10B
11B
9Be
10Be
(g,p)
(p,)
Extensive Nuclear Networks
 Automatic Adaptive Network
(n,g)
(g,n)
(g,)
(n,)
(,g)
b,(n,p)
Strong reactions
Weak reactions
Neutrinos
Initial stellar parameter
(mass, chemical composition)
Opacities
Equation of State
First model at the beginning
of the Pre-MS
Adaptive re-zoning
Atmosphere
Definition of
Convective
borders
MAIN PROGRAM
(Finite difference Henyey Method)
Mixing
Mass loss
Physical
evolution
THE FRANEC
CODE
Chemical
evolution
Output
New
temporal step
 AGB
 Thermonuclear SNe
 Core Collapse SNe
Evolution of Low & Intermediate Mass
Stars
Schematic structure of Schematic
an AGBstructure
star
of an AGB star
(not to scale)
(not to scale)
H-rich
convective
envelope
H-burning
shell
He-burning
shell
He
intershell
C-O core
Dredge-up
Flash-driven
intershell convection
Evolutionary track toward the WD
M=1 M
PN
0.6
CO
0.5 He
t =10 Gyr
Remnant:
CO WD
0.6 M
0.6
CO
AGB
0.55 He
0.2 CO
HB
RGB
0.1 He
WD
MS
Prada Moroni &
Straniero 2002
A WD in a binary system
toward a thermonuclear explosion
WD + 
2 WDs
“Universally” accepted model for Ia:
Thermonuclear Explosion
of a CO WD
M~MChandrasekhar ~ 1.4 M
Light Curve
56Ni

56Co
56 Fe
L
time
Supernova Cosmology Project
Lmax  MNi
WD is degenerate
e- Degenerate Pressure (EOS)
 1926 Fowler  Pauli Exclusion Principle
Pressure for relativistic electrons:

3 

1
2 3
PR
4
Z  
c

 A mH 
4
3
The Chandrasekhar limit
2
M Ch
 2 
1.456   M 
 e 
P independent of T


Thermonuclear
Explosion
tnuc < thyd
Thermonuclear Explosions
WD
WD
RG
SD
WD
DD
MCh
Propagation of the burning front
Detonation vburnvsound
Deflagration vburn< vsound
Delayed detonation
Deflagration Detonation
Compressional
heating
ignition
WD
C or He
detonation
C-deflagration
C-delayed detonation
Still Key Problems to control
SNIa !!
 Progenitors ? CO WD + companion
SD vs DD… both ??
Accretion ??
CSM: 2002ic
Hamuy et al. Nature 2003
2005gj Aldering et al. 2005
2006X Patat et al. Science 2007 NORMAL SNIa
 Explosion Mechanism ? begin subsonic
1D parametrization
3D still … fighting !!
(Barcelona, Chicago, MPI, NRL)
Massive 
Core
Collapse
At the end...
Layered Structure
Dense Iron Core
  107 g·cm-3
T  1010 K
MCore  1.4M
RSi-Core  4000 km
RFe-Core  800 km
Massive 
Core
Collapse
Fusing
H
He
C
Ne
O
Si
Main Fusion Products
He
C, O
Ne, O
O
S, Si, Ar
Fe, Cr
Time
6 million years
700000 years
1000 years
9 Months
4 Months
1 day
End result ?
A star whose core looks like an onion
Si Burning
54Fe, 56Fe, 55Fe,
58Ni, 53Mn
O Conv. Shell
28Si, 32S, 36Ar,
40Ca, 34S, 38Ar
C Conv. Shell
16O
28Si
“Fe”
24Mg,25Mg, 27Al
+ s-process
He Centrale
16O, 12C
He Shell
16O, 12C
H
Centrale+Shell
14N, 13C, 17O
H Centrale
4He
20Ne
20Ne, 23Na,
H Shell
He Shell
He Centrale
Main Products
C conv. Shell
Burning Site
O conv. Shell
Chieffi & Limongi
Si burning(Cent.+Sehll)
M=25M
12C
+ sprocess
Collapse and Explosion
1H
Core-Collapse Mechanism
Once the star has
finished its fuel the
core cools because of
two reasons:
a) Iron dissociation  fusion of light nuclei  the star
continues emitting energy
b) Degenerate e- gas  p + e-(2.25 MeV) n + e
(neutronization)  e escape and remove energy
c) Contraction turns into a free-fall collapse,
vast amount of neutrinos are produced
In less than 1 second the inner core radius goes from
4000 km to 10 km
(matter from the rest of the core is falling inward)
Core-Collapse Mechanism
Making Stars Explode
Because the neutrinos free path is small
the falling matter becames very hot
and expands outwards.
Finally, the star explodes and
ejects the star’s outer
layers into space.
All that remains of is a very
dense object: neutron star or black hole
PROBLEM: Turning the implosion into an explosion !!!
There are several models explaining the explosion,
but until now
simulations do not succeed in obtaining an explosion
Core Collapse SNe: LCs
Explosion Mechanism
Still Uncertain
simulated by a piston
of initial velocity v0,
located near the edge
of the Fe core
II-P
1. Rise: thermal energy
(envelope is fully ionized)
2. Plateau: recombination of H
Lenght  MH
3. Radioactive Tail: 56Co decay
L  M56Ni
56Ni  56Co 56 Fe
II-L No Plateau
Small H-envelope
Numerical Methods
 STELLAR EVOLUTION
FRANEC (Chieffi, Domínguez, Imbriani, Limongi, Piersanti, Straniero)
1D Hydrostatic Code
 Extended Nuclear Network (700 isotopes)
 Physics and Chemestry coupled
 Time dependent mixing
Low-mass 
PMS  AGB  WD 
 TPs
Massive 
PMS  Fe-core
Accretion  Explosive C-ignition
Numerical Methods
 EXPLOSION & LIGHT CURVES
1D Radiation-Hydrodynamic Code (PPM)
(Höflich, Khokhlov )
 Extended Nuclear Network (postprocess)
 Radiation transport via moments eq.
 Expansion opacities (scatt., bf, bb)
 Explosion mechanism: detonation  SNIa
deflagration
piston  CCSNe
 g Ray transport  Monte Carlo
LCs
 Frequency dependent transport eq. (1000 )
+
Eddington fac.
Mean opacities
Observations
LCs
Spectra (evolution)
Observed Relations
2001el SNIa
Krisciunas et al. 2003
1999em IIP
Hamuy et al. 2001
Lmax  LC
Lmax  B-V
Lmax VCa
Lmax VNi
1999ee SNIa
Hamuy et al. 2002
Information from the spectra
Hoflich et al. 2000
+ 15 days
Star of Si burning
-4 days
C-burning
MgII
SN1999by
1.05m
SNIa
Sub-L
CaII
1.15m
Duration of these phases  lower limit to the mass
SN
Remnants
Crab Nebula
SN 1054
Visible
X-ray
IR
Radio
Type Ia SN remnants:
shocked ejecta
O
Fe Si
S
Ca
Fe
Ar
XMM-Newton
Tycho
SN 1572
Interaction with the
Ambient Medium
AM~ 10-24 g/cm3
PDDT
DDT

X-ray emission spectra
  T Xi ionization
Ca
PDDT
Fe
Sub-Ch
Identify Explosion
Mechanism  DDT
Badenes et al. 2003
Cas A
Chandra
Si
Fe
Hwang et al. 2004
Age ~ 300 yr
SN1680
Good spatial resolution
X and Optical data
 CCSNe He-rich envelope
SiXIII/MgXI
Asymmetrically expanding
 Explosion ??
Vink et al. 2004
Bibliography
 BÖHM-VITENSE 1993, Introduction to Stellar
Astrophysiscs, University of Chicago Press.
 CLAYTON 1992, Principles of Stellar Evolution and
Nucleosynthesis, University of Chicago Press.
 HANSEN & KAWALER 1994, Stellar Interiors: Physical
Principles, Structure and Evolution, Springer-Verlag
 KIPPENHAHN 1990, Principles of Stellar Structure and
Evolution, Springer-Verlag.
 OSTLIE & CARROLL 1996, An Introduction to Modern
Stellar Astrophysics, Addison Wesley.
Bibliography
 PAGEL 1997, Nucleosynthesis and Chemical Evolution of
Galaxies, Cambridge University Press.
 BUSSO, GALLINO, WASSERBURG 1999,
Nucleosynthesis in AGB stars, Ann. Rev. A. &A., 36, 369.
 WALLERSTEIN et al. 1998, Synthesis of the elements
in stars forty years of progress,
Reviews of Modern Physics, Volume 69,