-
Nuclear AstrophysicsStellar evolution, core-collapse supernova
and explosive
nucleosynthesis
Karlheinz Langanke
GSI & TU Darmstadt & FIAS
Tokyo, December 2, 2009Karlheinz Langanke ( GSI & TU
Darmstadt & FIAS) Nuclear Astrophysics Tokyo, December 2, 2009
1 / 20
-
What is nuclear astrophysics?
Nuclear astrophysics aims at understanding the nuclear
processesthat take place in the universe. These nuclear processes
generateenergy in stars and contribute to the nucleosynthesis of
the elements.
3. The solar abundance distribution
+ +
Elemental(and isotopic)compositionof Galaxy at location of
solarsystem at the timeof it’s formation
solar abundances:
Bulge
Halo
Disk
Sun
Hydrogen mass fraction X = 0.71
Helium mass fraction Y = 0.28
Metallicity (mass fraction of everything else) Z = 0.019
Heavy Elements (beyond Nickel) mass fraction 4E-6
0 50 100 150 200 250mass number
10-13
10-12
10-11
10-10
10-9
10-8
10-7
10-6
10-5
10-4
10-3
10-2
10-1
100
num
be
r fra
ctio
n
D-nuclei12C,16O,20Ne,24Mg, …. 40Ca
Fe peak(width !)
s-process peaks (nucle ar shell closures)
r-process peaks (nucle ar shell closures)
AuFe Pb
U,Th
GapB,Be,Li
general trend; less heavy elements
N. Grevesse and A. J. Sauval, Space Science Reviews 85, 161
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 2 / 20
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Stellar reaction rateConsider Na and Nb particles per cubic
centimeter of particle types a and b.The rate of nuclear reactions
is given by:
r = NaNbσ(v)v
In stellar environment the velocity (energy) of particles
follows a thermaldistribution that depends on the type of
particles.
Nuclei (Maxwell-Boltzmann): φ(v) = N4πv2` m
2πkT
´3/2 exp “−mv22kT ”The product σv has to be averaged over the
velocity distribution φ(v)
〈σv〉 =∫ ∞
0
∫ ∞0
φ(va)φ(vb)σ(v)vdvadvb
Changing to center-of-mass coordinates, integrating over the
cm-velocity andusing E = µv2/2
〈σv〉 =(
8πµ
)1/2 1(kT )3/2
∫ ∞0
σ(E)E exp(− E
kT
)dE
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 3 / 20
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Charged-particle cross section
Stars’ interior is a plasma made of charged particles (nuclei,
electron).Nuclear reactions proceed by tunnel effect. For p + p
reaction Coulombbarrier 550 keV, but the typical energy in the sun
is only 1.35 keV.
cross section: σ(E) = 1E S(E)e−2πη; η = Z1Z2e
2
~
õ
2E =b
E1/2
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 4 / 20
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Astrophysical S factor
S factor allows accurate extrapolations to low energy.
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 5 / 20
-
Gamow window
Using definition of S factor:
〈σv〉 =(
8πµ
)1/2 1(kT )3/2
∫ ∞0
S(E) exp[− E
kT− b
E1/2
]dE
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 6 / 20
-
Gamow windowAssuming that S factor is constant over the Gamow
window andapproximating the integrand by a Gaussian one gets:
〈σv〉 =(
2µ
)1/2∆
(kT )3/2S(E0) exp
(−3E0
kT
)
E0 = 1.22[keV](Z 21 Z22µT
26 )
1/3
∆ = 0.749[keV](Z 21 Z22µT
56 )
1/6
(Tx measures the temperature in 10x K.)Examples for solar
conditions:
reaction E0 [keV] ∆/2 [keV] Imax T dependence of 〈σv〉p+p 5.9 3.2
1.1× 10−6 T 3.9
p+14N 26.5 6.8 1.8× 10−27 T 20α+12C 56.0 9.8 3.0× 10−57 T 42
16O+16O 237.0 20.2 6.2× 10−239 T 182
It depends very sensitively on temperature!Karlheinz Langanke (
GSI & TU Darmstadt & FIAS) Nuclear Astrophysics Tokyo,
December 2, 2009 7 / 20
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The solar pp chains
11H +
11H
21H +
11H
32He +
32He
21H + e
+ + νe
32He + γ
42He + 2
11H
32He +
42He
74Be + γ
85% 15%
(PP I)
74Be + e
− 73Li + νe
73Li +
11H 2
42He
(PP II)
74Be +
11H
85B + γ
85B
84Be + e
+ + νe
84Be 2
42He
(PP III)
15% 0.02%
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 8 / 20
-
The other hydrogen burning: CNO cycle
requires presence of 12C as catalyst
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 9 / 20
-
Energy generation: CNO cycle vs pp-chains
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 10 / 20
-
Consequences
stars slightly heavier than the Sun burn hydrogen via CNO
cyclethis goes significantly faster; such stars have much
shorterlifetimes
mass [M�] timescale [y]0.4 2× 10110.8 1.4× 10101.0 1× 10101.1 9×
1091.7 2.7× 1093.0 2.2× 1085.0 6× 1079.0 2× 10716.0 1× 10725.0 7×
10640.0 1× 106
hydrogen burning timescalesdepend strongly on mass.
Starsslightly heavier than the Sun burnhydrogen by CNO cycle.
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 11 / 20
-
Helium burning reactions
Critical Reactions in HeCritical Reactions in
He--burningburning
Oxygen-16Energy source in stellar He burningEnergy release
determined by associated reaction rates
Products of helium burning are carbon and oxygen, the bricks of
life!Karlheinz Langanke ( GSI & TU Darmstadt & FIAS)
Nuclear Astrophysics Tokyo, December 2, 2009 12 / 20
-
Carbon Burning
21
Burning conditions:
for stars > 8 Mo (solar masses) (ZAMS)
T~ 600-700 Mio
U ~ 105-106 g/cm3
Major reaction sequences:
dominates
by far
of course ps, ns, and as are recaptured 23Mg can b-decay into
23Na
Composition at the end of burning:
mainly 20Ne, 24Mg, with some 21,22Ne, 23Na, 24,25,26Mg,
26,27Al
of course 16O is still present in quantities comparable with
20Ne (not burning yet)
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 13 / 20
-
Evolution of massive star
Nuclear burning stages(e.g., 20 solar mass star)
28Si�J�D)0.023.5Ti, V, Cr,
Mn, Co, NiFeSi
16O + 16O0.82.0Cl, Ar,
K, CaSi, SO
20Ne�J�D)16O 20Ne�D�J)24Mg31.5Al, PO, MgNe
12C + 12C1030.8NaNe,
MgC
3 He4 Æ 12C12C�D�J)16O10
60.218O, 22Nes-process
O, CHe
CNO
4 H Æ 4He1070.0214NHeH
Main
Reaction
Time
(yr)
T
(109 K)Secondary
Product
Main
ProductFuel
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 14 / 20
-
Presupernova star
Star has an onion like structure.Iron is the final product of
thedifferent burning processes.As the mass of the iron coregrows it
becomes unstable andcollapses when it reaches around1.4 solar
masses.
0 50 100 150 200 250
0
1
2
3
4
5
6
7
8
9
A
Eb/A
1H
2H
6Li
4He
12C
16O
24Mg40Ca
56Fe 86Kr 107Ag 127I 174Yb 208Pb 238U
3He
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 15 / 20
-
A few facts: SN1987A
Type II supernova in LMC (∼ 55 kpc)
Egrav ≈ 1053 ergErad ≈ 8× 1049 ergEkin ≈ 1051 erg = 1 foe
neutrinos Eν ≈ 2.7× 1053 erg
0.0 5.0 10.0 15.0Time (seconds)
0.0
10.0
20.0
30.0
40.0
50.0
Ene
rgy
(MeV
)
Kamiokande IIIMB
light curve
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 16 / 20
-
Presupernova and collapse models
Core-collapse supernova simulations are separated into:1
presupernova models:
describes the stellar evolution through the various
hydrostaticburning stages (H, He,...,Si) and follows the collapse
of the centralcore until densities of order ρ9 = 10 are
reachedlarge nuclear networks are used to include the nuclear
energygeneration and the changes in compositionneutrinos, produced
in weak-interaction reactions, can leave thestar unhindered and are
treated as energy loss
2 collapse modelsdescribes the final collapse and the explosion
phasethe temperature during these phases is high enough that
allreactions mediated by the strong and electromagnetic
interactionare in equilibrium; thus the matter composition is given
by NuclearStatistical Equilibrium (NSE)reactions mediated by the
weak interaction are not in equilibriumneutrino interactions with
matter have to be considered in details(Boltzmann transport)
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 17 / 20
-
Core-collapse supernova.
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 18 / 20
-
Collapse phase
Important processes:T > 1010 K, ρ > 1010 g/cm3
Neutrino transport (Boltzmannequation):ν + A� ν + A (trapping)ν
+ e− � ν + e− (thermalization)
cross sections ∼ E2νelectron capture on nuclei andprotons:e−+
(N,Z ) � (N + 1,Z − 1) + νee− + p � n + νecapture on nuclei
dominates
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 19 / 20
-
Shell model and (d ,2He) GT strengths
C. Bäumer et al. PRC 68, 031303 (2003)
B(G
T+)
large shell modelcalculation
0.2
0.3
0.4B(G
T+ )
0 1 2 3 4 5 6 7Ex [MeV]
0.1
0.1
0.2
0.3
0.451V(d,2He)51Ti
0 2 4 6 8 10T (109 K)
10−8
10−7
10−6
10−5
10−4
10−3
10−2
λ (s
−1 )
LMP(d,2He)
0 2 4 6 8 10
0.6
0.8
1
1.2
1.4
Old (n, p) data
0 2 4 6 8 100
0.1
0.2
0.3
0.4
0.5
B(G
T+
)
51V(n,p) Alford et al. (1993)
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 20 / 20
-
Pauli blocking of Gamow-Teller transition
f7/2
p1/2
f5/2
p3/2
g9/2
N=40Blocked
� � � � � � � � � � � � �� � � � � � � � � � � � �� � � � � � �
� � � � � �
� � � � � � � � � � � �� � � � � � � � � � � �� � � � � � � � �
� � �neutrons protons
GT
Core
f7/2
p1/2
f5/2
p3/2
g9/2
CorrelationsFinite T
Unblocked
� � � � � � � � � � � � �� � � � � � � � � � � � �� � � � � � �
� � � � � �� � � � � � � � � � � � �
� � � � � � � � � � � �� � � � � � � � � � � �� � � � � � � � �
� � �� � � � � � � � � � � �
GT
neutrons protons
Core
Unblocking mechanism: correlations and finite
temperaturecalculation of rate in SMMC + RPA model
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 21 / 20
-
How do shell-model rates compare to previous rates?
1 4 7 10
T9
10-4
10-3
10-2
10-1
10-4
10-3
10-2
10-1
100
λ ec
(s−1
)
10-15
10-12
10-9
10-6
10-3
LMPFFN
1 4 7 10
T9
10-6
10-4
10-2
100
10-6
10-5
10-4
10-3
10-2
10-1
10-3
10-2
10-1
100
ρ7=10.7
56Fe
56Ni
ρ7=4.32
ρ7=4.32
55Co
59Co
ρ7=33
54Mn
ρ7=10.7 ρ7=33
60Co
A < 65: SM rates smaller
1010 1011 1012
ρ (g cm−3)
10−4
10−3
10−2
10−1
100
101
102
103
104
Rec
(s−
1 )
protonsnuclei
A > 65: SM rates larger
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 22 / 20
-
Neutrino trapping
ν + A � ν + A (trapping)elastic process, no energy, butmomentum
transferν + e− � ν ′ + e− (thermalization)inelastic scattering,
energy transferν + (Z ,A)→ ν ′ + (Z ,A)∗(thermalization)inelastic
scattering, energy transfercross sections ∼ E2ν
treatment by neutrino transport(Boltzmann equations) which
consider allneutrino types and keep track of neutrinofluxes,
energies at all space-time points
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 23 / 20
-
Effect on improved capture rates on collapse
With Rampp & Janka (General Relativistic model)15 M�
presupernova model from A. Heger & S. Woosley
1010 1011 1012 1013 1014
ρc (g cm−3)
0.20
0.25
0.30
0.35
0.40
0.45
Ye,
c, Y
lep,
c
BruennLMSH
Ylep
Ye
1010 1011 1012 1013 1014
ρc (g cm−3)
0.0
0.5
1.0
1.5
2.0
2.5
3.0
s c (
k B),
Tc
(MeV
)
BruennLMSH
T
s
1010 1011 1012 1013 1014
ρc (g cm−3)
10−1
100
101
102
103
104
105
106
Eν−
emis
sion
(M
eV s−
1 )
BruennLMSH
1010 1012 1014ρc (g cm−3)
0
10
20
30
40
50
〈Eν〉
(M
eV)
For ρ > 1012 g/cm3 fermi sea of neutrinos forms as neutrinos
gettrapped. Weak interaction is then basically in equilibrium!
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 24 / 20
-
Core bounceThe collapse continues until the centraldensity
becomes substantially (by about afactor 2-4) larger than nuclear
density(ρnm ≈ 2× 1014 g/cm3). Then nuclearpressure slows down the
infall and finallystops it. When the inner core has reached
itsmaximum density (maximum scrunch), itrebounds and a shock
starts.
A decisive quantity for this stage of thecollapse is the
Equation of State. It isassumed that matter consists of nuclear
andelectron components, while neutrinos havenegligible
interactions, but are important forthe determination of quantities
like Ye ortemperature.
In the shock the temperature increases. Sothe passage of the
shock dissociates thenuclei into free nucleons which costs theshock
energy (about 8-9 MeV/nucleon). Theshock has not enough energy to
traverse theiron core. It stalls. No prompt explosion.
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 25 / 20
-
Shock revival
The important reactions directlybehind the shock are:νe + n↔ p +
e−; ν̄e + p ↔ n + e+
Competition between emission(cooling) and absorption (heating)
byneutrinos.Thus the material directly behind theshock gets
heated.This increases the kinetic energy ofmatter and revives the
shock (delayedsupernova mechanism).However, spherical simulations
failand show no successful explosions.
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 26 / 20
-
Bounce and shock wave evolution
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 27 / 20
-
Bounce and shock wave evolution
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 27 / 20
-
Bounce and shock wave evolution
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 27 / 20
-
Bounce and shock wave evolution
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 27 / 20
-
Bounce and shock wave evolution
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 27 / 20
-
Bounce and shock wave evolution
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 27 / 20
-
Bounce and shock wave evolution
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 27 / 20
-
Bounce and shock wave evolution
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 27 / 20
-
Convection
There exist now two-dimensionalsimulations (with neutrino
transportand modern microphysics) whichyield successful
explosions.Convection brings neutrinos fromdeeper (hotter) layers
to the shockand increase the effectiveness ofenergy transfer.
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 28 / 20
-
Successful two-dimensional supernova
Successful 2-dimensional explosion of 11M� star with ONeMg
core(H.-Th. Janka)
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 29 / 20
-
Explosive nucleosynthesis in supernova
ννν
νν
ννν
νν
νν
νν
νShock Region
Explosive Nucleosynthesis
o
Consistent treatment of supernovadynamics coupled with a
nuclearnetwork.Essential neutrino reactions in theshock heated
region
νe + n � p + e−
ν̄e + p � n + e+
early (∼ 1 s): matter protonrich→ νp-processlater: matter
neutronrich→ r-process
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 30 / 20
-
The νp-process: basic idea
Protonrich matter isejected under theinfluence of
neutrinoreactionsNuclei form at distancewhere a
substantialantineutrino flux ispresent 0 10 20 30 40 50 60 70 80 90
100Mass Number A
10−310−210−1100101102103104105106107
Yi/Y
i,⊙
With νWithout ν
Antineutrinos help to bridge long waiting points via (n,p)
reactions
ν̄e + p → e+ + n; n + 64Ge→ 64Ga + p; 64Ga + p → 65Ge; . . .
C. Fröhlich, G. Martinez-Pinedo, et al., PRL 96 (2006)
142502
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 31 / 20
-
r-process
T ≈ 100 keV n & 1020 cm−3 implies τn � τβ(n, γ)� (γ, n)
implies Sn ≈ 2 MeV
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 32 / 20
-
The r-process at magic neutron numbers
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 33 / 20
-
Why are there r-process peaks?Once the path reaches nuclei with
magic neutron numbers (Z ,Nmag),the neutron separation energy for
the nucleus (Z ,Nmag + 1) decreasesstrongly. Thus, (γ,n) hinders
the process to continue and (Z ,Nmag)beta-decays to (Z + 1,Nmag −
1), which is followed immediately byn-capture to (Z + 1,Nmag). This
sequence of alternative β-decays andn-captures repeat itself, until
n-capture on a magic nucleus cancompete with destruction by
(γ,n).Thus, the r-process flow halts at the magic neutron numbers.
Due tothe extra binding energy of magic nuclei, the Qβ values of
these nucleiare usually smaller than those for other r-process
nuclei. This makesthe lifetimes of the magic nuclei longer than
lifetimes of other r-processnuclei. Furthermore, the lifetimes of
the magic nuclei increasesignificantly with decreasing neutron
excess. For example, the halfliveof the r-process nucleus 130Cd has
been measured as 195± 35 ms,while typical halflives along the
r-process are about 10 ms.Thus, material is enhanced in nuclei with
Nmag , which after freeze-out,results in the observed r-process
abundance peaks.
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 34 / 20
-
Where does the r-process occur in nature?
This has been named one of the 11 fundamental questions in
science.Recent observational evidence in metal-poor (very old)
stars point totwo distinct r-process sites. One site appears to
produce the r-processnuclides above A ∼ 130; another one has to add
to the abundance ofr-process nuclides below A = 130.The two
favorite sites are:
1 neutrino-driven wind above the proto-neutron star in
acore-collapse supernova
2 neutron star mergers
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 35 / 20
-
Which nuclear ingredients are needed?
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 36 / 20
-
Mass predictions
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 37 / 20
-
Half-lives for r-process nuclei
40 42 44 46 48 50Charge Number Z
10-3
10-2
10-1
100
T1/
2 (s
)
ETFSIFRDMHFBSMExpt.
N=82
66 67 68 69 70 71 72 73Charge Number Z
10-2
10-1
100
101
T1/
2 (s
)
ETFSIFRDMHFBSM
N=126
40 42 44 46 48 50Charge Number Z
0
20
40
60
80
100
Pβ,
n (%
)
FRDMSM
N=82
65 66 67 68 69 70 71 72 73Charge Number Z
0
20
40
60
80
100
Pβ,
n (%
)
FRDMSM
N=126
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 38 / 20
-
Supernova remnants
The remnant left over in the explosion depends on the
main-sequencemass Mms and on the maximum mass for neutron stars.
The later isnot quite well known. Most neutron stars, whose masses
are welldetermined (they are in binaries), have masses around 1.4
M�,however, recent observations might imply masses up to 2.1 M�.It
is generally assumed that the collapse of stars withMms > 20−
25M� leads to a black hole in the center, while stars with8M� <
Mms < 20− 25M� yield a supernova with a neutron starremnant.It
is also possible that accretion during the explosion might put
theremnant over the neutron star mass limit. It is speculated that
thishappened in the case of the SN87A.
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 39 / 20
-
Supernova remnants
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 40 / 20
-
Light curve
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 41 / 20
-
Energy from radioactive decays
A core-collapse supernova produces about 0.15− 0.2 M� 56Ni. This
ismade in the outer layers of the star (Ye = 0.5, mainly 16O) when
theshock wave passes through and brings this matter into NSE by
fastreactions. Supernova also produce other radioactive nuclides
(forexample 57Ni and 44Ti). 44Ti is only barely made (about 10−4
M�), buthas a lifetime of about 60 years. It dominates the
lightcurve of SN87Atoday.These radiactive nuclides decay, producing
γ radiation in the MeVrange. By scattering with electrons, these
photons are thermalized andthen radiated away as infrared, visible,
and ultraviolet light.
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 42 / 20
-
Radioactive decay
Light curve follows the decay of Nickel.
56Ni6 days−−−→ 56Co 77 days−−−−→ 56Fe
Tiempo
Núm
ero
de n
úcle
os
T1/2
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 43 / 20
-
Supernova neutrino detection
neutrino detection (SN1987A)
0.0 5.0 10.0 15.0Time (seconds)
0.0
10.0
20.0
30.0
40.0
50.0
Ene
rgy
(MeV
)
Kamiokande IIIMB
The Kamioka and IMB detectors are water Cerenkov
detectors.Observed have been ν̄e neutrinos via there interaction on
protons (inthe water molecule). The detection of the other neutrino
types is themain goal for the next nearby supernova to test the
predicted neutrinohierarchy.
Karlheinz Langanke ( GSI & TU Darmstadt & FIAS) Nuclear
Astrophysics Tokyo, December 2, 2009 44 / 20