neutrinos and neutrons in the r-process · most neutrinos emitted during the first ∼ 10 sec...

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