Supernova Nucleosynthesis and Neutrinos
Gail McLaughlin
North Carolina State University

Some Locations for Nucleosynthesis
Proto-neutron star supernovae• Explosive Burning yesterday’s talk by Heger• Neutrino winds originally discussed by Woosley et al, Meyer et al, Takahashi et al
Other core collapse supernovae, e.g. collapsars, hypernovae or similar• Explosive Burning e.g. Maeda & Nomoto, Nomoto & Tominga• Disk outflow, neutrino driven or otherwise Surman et al, Metzger et al.
• Jets e.g. Winteler et al. 2012, Nakamura et al. 2012,Papish et al. 2013
Compact Object Mergers• tidal tail ejection• Disk wind
• Disk disintegration

Types of Nucleosynthesis
Four types of interesting nucleosynthesis• r-process• p-process• 56 Ni and other iron group nuclei• 44 Ti see talk by Grebenev next
Necessary Physics• hydrodynamics• EOS, neutrino opacities• nuclear properties, reactions• neutrino flavor transformation

Interesting Questions• What is the site of the r-process?
• What the is the site of the p-process elements? LEPP?
• What nuclei do core collapse supernovae make?
• What nuclei do compact object mergers make?
• What nuclei do hypernovae make?
• What nuclei do disks (collapsar, mergers) make?

Avenues for exploration• What can be learned from astronomical observation?
• What can be learned from neutrino physics?
• What can be learned from supernova/neutron star models?
• What can be learned from a detailed analysis of the abundance
pattern?

Astronomical observation: r-process
1.00
0.50

SnCd

RuGa
Sr

0.00

log εPtZr
Nd Dy
GdCe Sm ErPbYbGe−0.50

OsBaPdYIrHfRh−1.00Pr EuNb−1.50

HoTb
CS 22892−052 Abundances
Upper Limits
SS r−Process Abundances

−2.00
−2.50

LaMo Ag

304050

AuTm
Lu60

ThU70

Atomic Number8090

Roederer et alCowan and Sneden

Halo (old) stars show a pattern above the second peak that is robust
and matches solar. e.g. Sneden et al, Honda et al., Frebel et al Second peak may be robust
as well. Roederer et al.

Inferences:
Inference 1: There is a different mechanism for the main r-process,for Ag, Pd (weak r-process?), and for N=50 nuclei (Sr, Y, Zr)
Inference 2: For the main r-process site, we want something that
happens early enough to allow for the population of a few halo stars,
and that is robust (produces the same pattern over and over again)
Traditional speculation: Core collapse supernovae are favored overneutron star mergers for the main r-process e.g. Argast et al, Wanajo et al

Most obvious site, supernova neutrino winds,
do not appear to make r-process
Need neutron rich material and the Ye in the wind is too high.
But could potentially get other things, like p-processSuggested mechanisms for p-process:• Ye ∼ 0.485 and some neutrino capture on nuclei Fuller and Meyer 1995
• νp -process - Ye slightly proton rich and antineutrino capture on
protons Frohlich et al
Back to r-process: basic inferences inconsistent without fixes - either
change the interpretation of the data or find a way to make SN workSome possible fixes• maybe supernovae form jets Winteler et al, Nakamura et al, Papish et al• maybe there are lots of “collapsars” around that make r-processfrom disk nucleosynthesis Surman et al• maybe lumps of material in 2D O-Ne-Mg supernovae maker-process Wanajo et al• maybe mergers evolve on a different timescale Korobkin et al• maybe halo stars come from dwarf galaxies without a lot of gas
around, i.e. maybe the stars are older than we think Frebel• maybe a sterile neutrino will help the neutrino wind
nucleosynthesis GCM et al, Tamborra et al.

What is required to make the r-process in mergers or supernovae
Two options:• eject material far from neutrinos and/or fast, and with littleheating• let the neutrinos set a low Ye (Eν̄e > Eνe )Why are the neutrinos important? They are in charge changing
interactions• ν̄e + p ↔ n + e+
• ν e + n ↔ p + e−Why don’t (active) neutrino oscillations
help the r-process in supernovae?
• oscillation increases the νe energy, makes νe + n → p + e− gofaster• oscillation tends to start late

Neutrino Flavor Transformation
Bipolar/nutation oscillations begin when the scale of the neutrino selfinteraction potential, Vνν ∼ GF (nν,ef f − nν̄,ef f ) approaches the scaleof the vacuum term δm2 /E.

• νe solid line• ν̄e dashed line
• green no oscillation• blue oscillation

Fig. from Duan et al. 2011

Small but non-negligible impact on nucleosynthesis
In SN winds, the oscillation often starts after nuclei begin to form

Early time density profile, s/k = 200, τ = 15ms

Late time density profile, s/k = 200, τ = 18ms

wind conditions tweaked to create r-process favorable conditions
For recent discussion of PNS winds, see Arcones and Thielemann 2012

Neutrino Flavor Transformation and p-processin supernovae
Abundance v. Mass Number
no oscillation
normal hierarchy
inverted hierarchy

100
10-1
10-2

abundance

10-3
10-4
10-5
10-6
10-7
10-8
10-9
10-10 0204060mass number

Fig. from C. Frohlich

Again, a small but non-negligible effect

80

100

120

Jets from Supernovae
Ejected Mass [M⊙ ]

10−2
10−3
10−4
10−5
10−6
10−760

80

100

120

140

160

180

200

220

240

Mass Number
Fig. from C. Winteler et al. 2012

If you have a magnetorotationally driven jet, it is possible to ejectmaterial with little heating and little influence of the neutrinos. e.g. Winteler
et al. 2012, Nakamura et al. 2012,Papish et al. 2013

Accretion Disk (Collapsar-type) Nucleosynthesis• in disks neutrino potential canbe negative• matter potential is positive• cancellation causes a new typeof neutrino transformation
red - no oscillations, blue - oscillationss/k = 50, β = 1.4 (moderate acceleration) figure from Malkus et al 2013

Neutrino Oscillations: scales

Modified wave equation

δm2
δm2ab
Ve + Vν − 4E cos(2θ)
Vν + 4E sin(2θ)
dψ
i~c ψν = 22δmadrVνb + δmsin(2θ)−V+−V+eν
4E4E cos(2θ)
Scales in the problem:•δm2vacuum scale 4E• matter scale Ve ∝ GF Ne (r)• neutrino self-interaction scale Vν ∝ GF Nν ∗ angle − GF Nν̄ ∗ angleVν has some subtleties. For proto-neutron star neutrinos Vν term
declines roughly as 1/r4

Neutrino Oscillations with Self-Interaction
Neutrino oscillations are with the self-interaction are numerically moredemanding than standard matter enhanced neutrino effects. Calculations see
Balantekin, Dighe, Duan, Fuller, Kajino, Kneller, Raffelt, Volpe, Yoshida and more

1. Non-linear effects Vν ∼ |ψν |2
2. Requires multi-group treatment: Background involves sum overneutrino momentum and angle.
We rewrite the problem as an S-matrix problem
dS= HSi~dx
and solve in an adiabatic basis. Kneller and McLaughlin

Accretion disk neutrino transition regions
Type I - Matter enhanced region• Traditional MSW region• vacuum interaction strength is the same size as matter potential• neutrino self interaction strength is small√
2• i.e δmij /Eν ∼ 2GF Ne ≫ VννType II - nutation/bipolar• “Traditional” nutation in NFIS picture (also called bipolar)• δm2ij /Eν ∼ Vνν• occurs closer to PNS than Type I regions• occurs when matter potential is both large and small

Accretion disk transition regions
Type III “Neutrino - Matter enhanced” New!
• Self-interaction potential ∼ matter potential• vaccuum interaction strength is small√
2• i.e. δmij /Eν << 2GF Ne ∼ Vνν• potentials have opposite signs! cancellation!
• occurs in both hierarchies• not a usual situation in supernovae

Massive star accretion disks:ν̄ dominated first, ν later
Malkus et al. 2012Survival Probability

1(III)

(III)

(I)

0.8
0.6
0.4(I)

0.2
0Interaction Strength (erg)-15

1×10-18

1×10-21

1×10-24

1×10

5

1×10

6

1×10

7

8

1×10
1×10
Progress (cm)91×10Upper panel: solid red - electron neutrino survival probability
Upper panel: dashed red - electron antineutrino survival probability

Flavor Transformation: PNS vs. Disks
Supernovae Neutrinos• In the SN, oscillations tend to occur after the most important
point for wind nucleosynthesis• In the SN, oscillations increase νe , ν̄e capture rates• There is some re-arrangement of the abundance pattern
Accretion disk neutrinos• In the accretion disk, they occur earlier, because of the“matter-neutrino enhanced” transition Malkus et al 2012
• In the “collapsar-type” accretion disk, oscillations decrease νe , ν̄ecapture rates• One expects significant changes the abundance patternAccretion Disk (Collapsar-type) Nucleosynthesis• in disks neutrino potential isnegative• matter potential is positive• cancellation causes a new typeof neutrino transformation
red - no oscillations, blue - oscillationss/k = 50, β = 1.4 (moderate acceleration) figure from Malkus et al 2013

Back to ideas for the r-process:
jets, disks, compact object mergers
They all claim to have enough neutrons...
How can we decide? Can radioactive beams help?

The rare earth peakSolar abundance data with the rare earth peak in red

Constaining conditions using abundance patterns
e.g. rare earth peak
• Determine how the rare earth peak forms• Consider the nuclear physics uncertainties• Determine the range of thermodynamic conditions that aresuccessful

How to read r-process flow plots

constant beta decay rate

stable isotopes

beta decayZ

constant neutron
photo diss.

n−captseparation energy(controls most populatedisotope)

most abundantisotopesNHow to form structures

ZN

Hot rare earth peak formation Mumpower, GCM, Surman 2012, see also Surman and Engel 1997

Neutron capture just below the peak and photo-dissociation above FRDM

Correct structure must persist until neutrons are exhausted
1553570
35
1570
5151551150.05

Calculation with different mass model ETFSI

0.20.5Reduce nuclear physics uncertainties:
Most important nuclei for the formation of the rare earth peak
Mumpower, GCM, Surman 2012

Conclusions• Exciting era in supernova nucleosynthesis• Calculations are improving as are observations• r-process not yet settled: compact object mergers, SN jets,
accretion disks• Neutrino oscillations are a significant factor in accretion disk
nucloesynthesis• Nailing down nuclear properties will give us additional clues toastrophysical conditions for example, that make the rare earth peakConstraints on the rare earth peak from solar data✭ ✁✭❜✁

Red:
comparison to data withGREP < 0.15
Yellow: rare earth to A=195 ratio towithin a factor of 2
Green: measures late time neutron
capture on the sides of the peak
∆GA=150−180
f inal/min < 1 see also Arcones and MartinezPinedo 2011✭✂✁Constrain cooling time scale T ∼ t−nand entropy/baryon s✭✄✁Calculation with the FRDM (finite range droplet model) using Ye = 0.3, Ye = 0.4