neutrinos and neutrons in the r-process · most neutrinos emitted during the first ∼ 10 sec...
TRANSCRIPT
Neutrinos and Neutrons
in the r-process
Gail McLaughlin
North Carolina State University
1
Neutrino Astrophysics
• What are the fundamental properties of neutrinos?
• What do they do in astrophysical environments?
• What do neutrinos in a core collapse supernova do?
• What do neutrinos in a black hole accretion disk do?
Nuclear Astrophysics
• What is the origin of the elements?
• Where are the heaviest elements made?
• What elements are made in compact object mergers/ stellar
explosions?
2
Where is uranium made?
Uranium is an r-process element . . .
3
From where do the
heaviest elements come?
• Elements beyond the iron peak in three main categories
– r-process
– s-process
– p-process
• examples of r-process elements
– Europium
– Platinum
– Uranium
– Thorium
4
Solar Abundances
1e-05
0.0001
0.001
0.01
0.1
40 60 80 100 120 140 160 180 200
Abu
ndan
ce
A
scaled solar data
5
The r-process elements
e. g. Uranium-238 Z=92, N=146 → need lots of neutrons
A(Z, N) + n ↔ A + 1(Z, N + 1) + γ
A(Z, N) → A(Z + 1, N − 1) + e− + ν̄e
Z
N
rapid neutron capture as compared with beta decay
At what temperatures does this happen?
(n, γ) (γ, n) equilibrium at T ≈ 109K
6
What astrophysical sites have a lot of neutrons
and eject material?
• neutron star mergers/ neutron star - black hole mergers
• “wind” close to the center of a core-collapse supernova
7
Two kinds of r-process data:
Observational and Meteoritic
Meteoritic: Isotopic measurements of r-process nuclei:
two r-process sites? Wasserburg, Busso and Gallino (1996)
Observational Halo Stars:
two r-process sites?
Figure from Cowan and Sneden (2004)
30 40 50 60 70 80 90Atomic Number
−2.50
−2.00
−1.50
−1.00
−0.50
0.00
0.50
1.00
log
ε
CS 22892−052 AbundancesUpper LimitsSS r−Process Abundances
U
Th
Pb
Os
Pt
Ce
Nd
Sm
Dy
La
Ba
Pr EuTb
Sr
Y
Zr
Ru
Nb
Mo Ag
Pd
Ho
Tm
Hf
Yb
ErGd
Ge
CdGa
Au
Ir
Lu
Rh
Sn
What do the data suggest?
Supernovae are favored over neutron star mergers Argast et al 2004
8
Core Collapse Supernovae
end of the life of a massive star
HeH
Oxygen
Si
Fe core
C
• core unstable
Mcore ∼ 1.5Msun
• collapse to nuclear density
• core bounce
• shock produced
• shock stalls
• neutrinos diffuse out of core,
may energize shock
9
Supernova Explosion!
Figure from John Blondin
10
Explosions of Massive Stars:
What’s happening at the Center?
proto neutronstar core
shockoutflow
black hole, jet
neutrinos
accretion disk
Standard core core collapse SNLong duration gamma ray burst
from core collapse SN or compact
object merger
11
Explosions of Massive Stars:
Where is the nuclear-neutrino physics?
proto neutronstar core
oscillationsneutrino
nucleosynthesis
of core
nuclear physics
neutrino
& emissionscattering
shockoutflow
neutrinos
accretion disk
black hole, jet
nucleosynthesis
neutrino oscillations
nuclear physicsof the disk
neutrino
& emission
scattering
Standard core core collapse SNLong duration gamma ray burst
from core collapse SN or compact
object merger
12
Neutrinos from Standard Supernovae
All types of neutrinos are trapped in the core. At the surface they
escape and travel through the outerlayers of the SN, then to earth.
νe
ν e
short mean freepath
long mean freepath
core
ν ν ν ν µ τ τ µ
The Neutrino Sphere, near
the surface of the proto-
neutron star is where the
neutrinos decouple. This is
where the neutrino energies
are determined.
13
Measuring the Supernova Neutrino Signal
Why? Neutrinos are our window deep into the core of the Supernova
Roughly the range is
• 〈Eνµ〉 = 〈Eν̄µ
〉 = 〈Eντ〉
= 〈Eν̄τ〉 = 16 − 25 MeV
• 〈Eν̄e〉 = 12 − 18 MeV
• 〈Eνe〉 = 8 − 13 MeV
Predictions of the spectral shape
are different, too.
14
Supernova Neutrinos
Most neutrinos emitted during the first ∼ 10 sec
Galactic supernovae estimated to occur ∼ 1 every 30 years
Supernova neutrinos detected from SN1987a:
∼ 20 events observed in Kamiokande and IMB.
15
Returning to the uranium problem
We want to examine a supernova “wind” near the core
We want to know if it has a lot of neutrons . . .
16
The weak interaction
The only way to convert protons to neutrons and vice-versa
• beta decay n → p + e− + ν̄e
• electron (neutrino) capture e− + p ↔ n + νe
• positron (antineutrino) capture e+ + n ↔ p + ν̄e
A few words about neutrinos . . .
They come in three flavors: νe, νµ, ντ and have associated charged
leptons e, µ, τ
The neutrinos have 10s of MeV energies, not enough to make muons
or taus, so only the electron type change charge.
17
The weak interaction
So does the material have more protons or neutrons?
• electron (neutrino) capture e− + p ↔ n + νe
• positron (antineutrino) capture e+ + n ↔ p + ν̄e
Electons and positrons are in equilibrium with the rest of the matter,
the neutrinos are not.
i.e. the electrons are positrons are cooling as material moves away
from the core, but the neutrino temperatures are fixed.
At some point the neutrino and antineutrino reactions begin to
dominate
18
The weak interaction
So does the material have more protons or neutrons?
• electron (neutrino) capture e− + p ↔ n + νe
• positron (antineutrino) capture e+ + n ↔ p + ν̄e
19
Let’s try a network calculation
Reaction network calculation contains ∼ 3000 elements and isotopes
Takes as input temperature, density profiles . . .
Outputs an abundance pattern . . .
Recall that we want to match the data . . .
1e-05
0.0001
0.001
0.01
0.1
40 60 80 100 120 140 160 180 200
Abu
ndan
ce
A
scaled solar data
20
Core Collapse Supernovae: Nucleosynthesis
in the Traditional Neutrino Driven Wind
Hoped for r-process site
1e-10
1e-09
1e-08
1e-07
1e-06
1e-05
0.0001
0.001
0.01
0.1
40 60 80 100 120 140 160 180 200
Abu
ndan
ce
A
scaled solar dataν-driven wind
What happened?
Not enough entropy
Eν̄es too low
Most current simulations
don’t work!
21
Nucleosynthesis in Hot Outflows
n,p →4 He → iron peak nuclei → heavier nuclei
22
Nucleosynthesis in hot outflows
What matters?
• outflow timescale, milliseconds to seconds
• entropy s ∼ 20 to s ∼ 400
• electron fraction, Ye = p
n+p, Ye ∼ 0.1 to 0.6,
Electron fraction is set by the weak interactions:
νe + n ↔ p + e−, ν̄e + p ↔ n + e+
What about ν̄e + p → n + e+ on left over protons?
23
Neutrino induced processes
in proton rich winds in core collapse supernova
Antineutrino capture on left over
protons produces p process ele-
ments
10-8
10-6
10-4
10-2
100
102
104
106
108
50 60 70 80 90 100 110 120 130 140 150A
BU
ND
AN
CE
S Y
/Ysu
n
MASS NUMBER A
standard(n,p)* 5.0(n,p)* 0.2
Frohlich et al 2011
24
But what about the r-process?
25
Core Collapse Supernovae: Nucleosynthesis
in non-Traditional Neutrino Driven Winds
1e-10
1e-09
1e-08
1e-07
1e-06
1e-05
0.0001
0.001
0.01
0.1
40 60 80 100 120 140 160 180 200
Abu
ndan
ce
A
scaled solar dataνe ⇔ νs oscillations
Active Sterile ν oscillations
Beun et al 2007
Other ideas:
• magnetic fields
• 3-D calculations
• compact object mergers
26
Compact Object Mergers:
Can they make the r-process?
• two neutrons stars spiral together
• or: a black hole and a neutron star spiral together
• a disk is formed
• some material is ejected from tails
• some material is blown off in a neutrino wind
Examine: neutrino wind from the disk
27
Neutrino Surfaces
Black Hole Neutron Star Merger
Accretion disk in neutrinos Side view shows neutrino surfacesSurman et al. 2008
28
Accretion Flow Nucleosynthesis
Black Hole Neutron Star Merger
Hypothetical schematic of events
in outflow
Do you get to the r-process stage
in this environment?
29
Uncertainties
• Astrophysical Environment
• Equation of state
• Neutrino Properties/Oscillations
• β-decay rates
• nuclear masses/capture rates
• fission rates/daughter products
30
Fission Cycling in the r-process
30 40 50 60 70 80 90 100−5
−4
−3
−2
−1
0
1Lo
g A
bund
ance
Atomic Number
CS 22892−052SS r−process Abundancesν
e ↔ ν
s oscillations
astrophysical data hints at
fission cycling
10−2
10−1
100
0
0.5
1
1.5
2
2.5
Ye
Y19
5 ÷ Y
130
Symmetric FissionAsymmetric FissionAsymmetric Fission + 2 n
fission cycling requires many
neutrons
Very little data on the relevant fission rates and daughter products
Beun et al 2006, Beun et al 2008
31
Neutron Capture rates on the A=130 peak
Change the mass model
(bottom) or vary the capture
rates (top)
48
49
50
52
80 83 86 87
126Cd Cd Cd Cd Cd Cd Cd Cd Cd
InInInInInInInInInIn
Sn
127
128
Cd
Sn Sn Sn Sn Sn Sn Sn Sn
Sb Sb Sb Sb Sb Sb Sb Sb Sb Sb
Te Te Te Te Te Te Te Te Te Te130
7978 82 84 85
130
130
130
130
81
129
129
129
129
128
128127
131 132 133 136 137 138 139134
131
131
131
131
132
132
132
132
133
133
133
133
134
134
134
134
135
135
135
135
136
136
136
137
137
13851
N
Sn
135
Z
Nuclei in the 130 peak with
neutron capture rates that
effect a 5% or more change in
the abundance distribution
Surman et al 2008
32
Neutron capture rates and the rare earth peak
150 160 170 180
10−5
10−4
10−3
A
Log[
Y(A
)]
156Nd160Pm168Gd165Pm167EuBaseline
Mumpower et al in prep 2011
Rare earth peak forms from combination of beta decay and neutron
capture Surman and Engel 1997
Neutron capture rates together with astrophysical conditions determine
the rare earth peak. Entering precision era of abundance measurements
→ need better data and better analysis of astrophysical conditions.
33
Summary
• Supernova may be the site of the r-process (or not)
• Neutrinos play an important role in determining the type of
elements formed in supernovae and compact object mergers
• Neutron capture as well as beta decay are important for future
understanding the r-process
• Improvements in our understanding of nuclear physics will allow
better constraints on the site of the r-process
34