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
0.1
scaled solar data
Abundance
0.01
0.001
0.0001
1e-05
40
60
80
100 120 140 160 180 200
A
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
Z
A(Z, N ) → A(Z + 1, N − 1) + e− + ν̄e
N
rapid neutron capture as compared with beta decay
At what temperatures does this happen?
(n, γ) (γ, n) equilibrium at T ≈ 109 K
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)
1.00
0.50
Sr
0.00
log ε
Sn
Cd
Ru
Ga
Observational Halo Stars:
two r-process sites?
Pt
Zr
Nd Dy
Gd
Ce Sm Er
Pb
Yb
Ge
−0.50
Os
Ba
Pd
Y
Ir
Hf
Rh
−1.00
La
Mo Ag
Nb
−1.50
Ho
Pr Eu
Tb
Tm
Figure from Cowan and Sneden (2004)
CS 22892−052 Abundances
Upper Limits
SS r−Process Abundances
−2.00
−2.50
30
40
50
60
Au
Th
Lu
U
70
80
Atomic Number
What do the data suggest?
Supernovae are favored over neutron star mergers
8
Argast et al 2004
90
Core Collapse Supernovae
• core unstable
Mcore ∼ 1.5Msun
end of the life of a massive star
Oxygen
Si
• collapse to nuclear density
C
He
• core bounce
H
• shock produced
Fe core
• 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?
black hole, jet
shock
outflow
neutrinos
accretion disk
proto neutron
star core
Standard core core collapse SN
Long duration gamma ray burst
from core collapse SN or compact
object merger
11
Explosions of Massive Stars:
Where is the nuclear-neutrino physics?
nucleosynthesis
neutrino
oscillations
black hole, jet
shock
nucleosynthesis
outflow
neutrino oscillations
neutrinos
accretion disk
nuclear physics
of core
neutrino
scattering
proto neutron
& emission
star core
Standard core core collapse SN
nuclear physics
of the disk
neutrino
scattering
& emission
Long 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.
long mean free
path
νe
core
νe
νµ νµ ντ ντ
short mean free
path
The Neutrino Sphere, near
the surface of the protoneutron 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
• hEνµ i = hEν̄µ i = hEντ i
= hEν̄τ i = 16 − 25 MeV
• hEν̄e i = 12 − 18 MeV
• hEνe i = 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 . . .
0.1
scaled solar data
Abundance
0.01
0.001
0.0001
1e-05
40
60
80
100 120 140 160 180 200
A
20
Core Collapse Supernovae: Nucleosynthesis
in the Traditional Neutrino Driven Wind
Hoped for r-process site
0.1
scaled solar data
ν-driven wind
0.01
Abundance
0.001
0.0001
What happened?
1e-05
Not enough entropy
1e-06
Eν̄e s too low
1e-07
Most current simulations
don’t work!
1e-08
1e-09
1e-10
40
60
80
100 120 140 160 180 200
A
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
108
Antineutrino capture on left over
protons produces p process elements
ABUNDANCES Y/Ysun
106
104
102
100
10-2
10-4
10-6
standard
(n,p)* 5.0
(n,p)* 0.2
-8
10
50
60
70
80
90
100 110
MASS NUMBER A
Frohlich et al 2011
24
120
130
140
150
But what about the r-process?
25
Core Collapse Supernovae: Nucleosynthesis
in non-Traditional Neutrino Driven Winds
0.1
scaled solar data
νe ⇔ νs oscillations
0.01
0.001
Abundance
0.0001
Other ideas:
1e-05
1e-06
• magnetic fields
1e-07
• 3-D calculations
1e-08
1e-09
1e-10
40
60
80
100 120 140 160 180 200
• compact object mergers
A
Active Sterile ν oscillations
Beun et al 2007
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 surfaces
Surman et al. 2008
28
Accretion Flow Nucleosynthesis
Black Hole Neutron Star Merger
Do you get to the r-process stage
in this environment?
Hypothetical schematic of events
in outflow
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
2.5
1
CS 22892−052
SS r−process Abundances
νe ↔ νs oscillations
2
1.5
÷Y
130
−1
195
−2
Y
Log Abundance
0
Symmetric Fission
Asymmetric Fission
Asymmetric Fission + 2 n
1
−3
0.5
−4
−5
30
40
50
60
70
80
Atomic Number
90
100
astrophysical data hints at
fission cycling
0 0
10
−1
10
Y
−2
10
e
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
52
130
131
132
133
134
135
136
137
138
139
51
129
130
131
132
133
134
135
136
137
138
Z 50
128
129
130
131
132
133
134
135
136
137
49
127
48
126
Te
Sb
Sn
In
Cd
78
Te
Sb
Sn
128
In
127
Cd
79
Te
Sb
Sn
129
In
128
Cd
80
Te
Sb
Sn
130
In
129
Cd
81
Te
Sb
Sn
Te
Sb
Sn
131
132
In
130
Cd
In
131
Cd
82
83
Te
Sb
Sn
133
In
132
Cd
84
Te
Sb
Sn
134
In
133
Cd
85
Te
Sb
Sn
135
In
134
Cd
86
Te
Sb
Sn
136
In
135
Cd
87
N
Change the mass model
(bottom) or vary the capture
rates (top)
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
−3
10
156
Nd
Pm
168
Gd
165
Pm
167
Eu
Baseline
Log[Y(A)]
160
−4
10
−5
10
150
160
170
180
A
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