[PPT] - Core Collapse Supernovae: Explosion models and long-term neutrino PowerPoint Presentation

SLIDE 1

2 1 1 2

Core Collapse Supernovae: Explosion models and long-term neutrino emission

Luke Roberts NSCL, MSU

SLIDE 2

Energy (MeV) Time (s) See Bionta et al. ’87 and Hirata et al. ‘87

Core Collapse Supernovae: Multi- Messenger Events

2

Neutrinos Gravitational Waves

From Ott et al. ‘12

Electromagnetic

From Filppenko ‘97

Nucleosynthesis

From Amarsi et al. ‘15

−0.5 0.0 0.5 1.0 1.5 −0.5 0.0 0.5 1.0 1.5 [O/Fe]

[OI] 630nm OI 777nm

− − − − − 1.0 − − −3 −2 −1 [Fe/H] −1.0 − [O/Fe] −

SLIDE 3

Overview

3D Central Engine Models
Long term CCSN neutrino emission

SLIDE 4

Core Collapse

GMNS

2

RNS ~ 3×1053 erg

▪ Stars with M >~ 9 Msun burn their core to Fe ▪ Core exceeds a Chandrasekhar mass supersonic collapse outside of homologous core bounce shock after ~2 x saturation density ▪ Gravitational binding energy of compact remnant:

~ 1051erg

▪ Binding energy of stellar envelope:

10 20 30 40

Ln (1052 erg s−1) −100 −50

50 100 150 200

Time Post Bounce (ms)

6 8 10 12 14 16

en (MeV)

SLIDE 5

Self Consistent Spherically Symmetric CCSN Explosions

5

1 2 3 4 L [10

52 erg s

1]

10

2

10

1

10 0.05 0.1 0.15 0.2 8 10 12 <ε> [MeV] 2 4 6 8 5 10

νe νe νµ/τ L/10

Time after bounce [s]

Accretion Phase Cooling Phase

0.3 0.5 1 2 3 5 10 20 30 50

Energy Luminosity [1051 erg/s]

ν

e

¯ ν

e

ν

µ /τ

7 Time After Bounce [s] 0.01 0.02 0.05 0.1 0.2 0.3 0.5 1 2 3 5 7 7 8 9 10 11 12 13 14 15 16 17 18 19 Time After Bounce [s]

Mean Energy [MeV]

ν

e

¯ ν

e

ν

µ /τ

¯ ν

µ /τ

Huedepohl et al. (2010) Fischer et al. (2010, 2012) Only possible for low mass progenitors, mainly ECSN

SLIDE 6

Simulating CCSNe

6

Hydrodynamics + General Relativity + Neutrino Transport + Microphysics (EoS, ν-opacities, nuclear network)

SLIDE 7

Post Bounce Evolution of CCSNe

Hydrodynamic instabilities (such as

convection and SASI) can aid energy transport and shock propagation

In axial symmetry, this enhances the

efficacy of neutrino energy deposition and results in successful explosions (Mueller et al. ’12, Bruenn et al. ’13)

Does the neutrino mechanism work

in 3D?

How does this depend on input

physics and numerics?

SLIDE 8

Two Moment Neutrino Transport

Entropy Take angular moments of the neutrino distribution function:

∂t ˜ E + ∂j ⇣ α ˜ Fj − βj ˜ E ⌘

+ ∂ν

⇣ ναnα ˜ Mαβγuγ;β ⌘

= α

h ˜ PijKij − ˜ Fj∂j ln α − ˜ Sαnα i ∂t ˜ Fi + ∂j ⇣ α ˜ Pj

i − βj ˜

Fi ⌘

− ∂ν

⇣ ναγiα ˜ Mαβγuγ;β ⌘

= α

 ˜ Fk∂iβk α

− ˜

E∂i ln α + ˜ Pjk 2 ∂iγjk + ˜ Sαγiα

⇣

⌘ h i

MAk

(ν)

=

Z

dVp pα1...pαk

(pµuµ)k2 f (pβ, xβ)δ(ν + pδuδ)

Z

Pαβ

(ν) = 3χ(ξ) − 1

2 Pαβ

(ν),thin + 3(1 − χ(ξ))

2 Pαβ

(ν),thick.

Get conservation equations for projections of the rest frame energy dependent stress tensor: Still need to specify neutrino stress tensor:

Amenable to finite volume techniques and truly 3D, but

Mαβ

(ν) ;β

∂xα ∂τ ∂ f (xµ, pµ) ∂xα

+ ∂pi

∂τ ∂ f (xµ, pµ) ∂pi

= ˜

S(xµ, pµ)

Boltzmann Equation:

See e.g., Shibata et al. ’11, Cardall et al. ’13, Just et al. ’15, Kuroda et al. ’16, LR et al. ‘16

SLIDE 9

Evolution to Explosion

9 2 4 6

L [1052 erg s−1]

e ¯ e x

40 80 120 160 200 240 280 320 360

t − tb [ms]

8 10 12 14 16 18 20

[MeV]

LR et al. (2016)

SLIDE 10

Resolution and Symmetry Dependence of CCSNe Models

10

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL 1

˙ M [M s1]

LR et al. (2016)

SLIDE 11

Resolution and Symmetry Dependence of CCSNe Models

11

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL 1

˙ M [M s1]

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL s27FH 1

˙ M [M s1]

LR et al. (2016)

SLIDE 12

Resolution and Symmetry Dependence of CCSNe Models

12

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL 1

˙ M [M s1]

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL s27FH 1

˙ M [M s1]

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL s27FH s27OH 1

˙ M [M s1]

LR et al. (2016)

SLIDE 13

Resolution and Symmetry Dependence of CCSNe Models

13

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL 1

˙ M [M s1]

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL s27FH 1

˙ M [M s1]

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL s27FH s27OH s27OL 1

˙ M [M s1]

LR et al. (2016)

SLIDE 14

Resolution and Symmetry Dependence of CCSNe Models

14

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL 1

˙ M [M s1]

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL s27FH 1

˙ M [M s1]

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL s27FH s27OH s27OL 1

˙ M [M s1]

40 80 120 160 200 240 280 320 360

t tb [ms]

100 150 200 250 300 350 400

Shock Radius [km]

s27FL s27FH s27OH s27OL 1

˙ M [M s1]

LR et al. (2016)

SLIDE 15

Turbulent Convection

15

50 100 150 200 250 300

Radius [km]

0.00 0.05 0.10 0.15 0.20 0.25 0.30 0.35

TrRij/p0/3

81 ms 128 ms 179 ms 228 ms 279 ms

Reynolds stress can contribute significantly to the pressure in the gain region and there is some resolution dependence of the Reynolds stress

Murphy & Meakin ’11, Handy et al. ’14, Couch & Ott ‘15

LR et al. (2016)

SLIDE 16

3D Explosion Models

16

d f D =

m

q º f q D = D = = 2 q D =

q

D =

n

n n n n =

mt m t

n n n =

mt m t

a =

r >
r <
a

r >

+

+ + + + +

+

Lentz et al. (2015) Janka et al. (2016)

Many groups are seeing shock runaway, but maybe not quantitative

agreement

Sensitive to input physics (Melson et al. ’15) and resolution (Radice et al. ’15)
Nevertheless, things look relatively positive for 3D shock runaway

Takiwaki et al. ’12, Melson ’15, Lentz ’15, LR et al. ’16, Takiwaki et al. ‘16

SLIDE 17

Jet Driven Supernovae

Rapidly rotating,

magnetized SNe

Full 3D Dynamics

also important here

Kink instabilities in

jet significantly change dynamics

17

= = = t tb = 186.4ms

Full 3D

t tb = 67.8ms

Octant symmetry

= = =

Figure 1. −

Moesta et al. (2016)

SLIDE 18

The Supernova Neutrino Signal

Super-Kamiokande Neutrino Detector ~20 Neutrino Events Observed from SN 1987a at two detectors via the reaction Larger, modern detectors will detect thousands of events from a nearby supernova, allowing us to directly probe the nature of the nascent neutron star

ν

e + p → e+ + n

See Bionta et al. ’87 and Hirata et al. ‘87

Energy (MeV) Time (s)

SLIDE 19

Milky Way Supernova Rate

19

Most recent known MW CCSN

Cas A (~300 yrs)

Look for supernovae in

galaxies analogous to MW (Cappellaro et al. 1999)

Take census of historical

galactic supernovae and correct for obscuration (Tammann et al. 1994)

Reasonably consistent

search galaxy rate [SNu] type Ia II+Ib/c All S0a-Sb 0.27 ± 0.08 0.63 ± 0.24 0.91 ± 0.26 Sbc-Sd 0.24 ± 0.10 0.86 ± 0.31 1.10 ± 0.32 Spirals∗ 0.25 ± 0.09 0.76 ± 0.27 1.01 ± 0.29

∗ Includes types from Sm, irregulars and peculiars.

Cappellaro et al. (1999)

multiply by ~2.4 to get MW rate

SLIDE 20

SN* Neutrino Detectors

20

Detector Type Mass (kt) Location Events Flavors Status Super-Kamiokande H2O 32 Japan 7,000 ¯ νe Running LVD CnH2n 1 Italy 300 ¯ νe Running KamLAND CnH2n 1 Japan 300 ¯ νe Running Borexino CnH2n 0.3 Italy 100 ¯ νe Running IceCube Long string (600) South Pole (106) ¯ νe Running Baksan CnH2n 0.33 Russia 50 ¯ νe Running MiniBooNE∗ CnH2n 0.7 USA 200 ¯ νe (Running) HALO Pb 0.08 Canada 30 νe, νx Running Daya Bay CnH2n 0.33 China 100 ¯ νe Running NOνA∗ CnH2n 15 USA 4,000 ¯ νe Turning on SNO+ CnH2n 0.8 Canada 300 ¯ νe Near future MicroBooNE∗ Ar 0.17 USA 17 νe Near future DUNE Ar 34 USA 3,000 νe Proposed Hyper-Kamiokande H2O 560 Japan 110,000 ¯ νe Proposed JUNO CnH2n 20 China 6000 ¯ νe Proposed RENO-50 CnH2n 18 Korea 5400 ¯ νe Proposed LENA CnH2n 50 Europe 15,000 ¯ νe Proposed PINGU Long string (600) South Pole (106) ¯ νe Proposed

Scholberg et al. (2015)

SLIDE 21

10−1 100 101 102

tpost-bounce (s)

4 6 8 10 12 14 16

en (MeV)

102 10−1 100 101 102 10−4 10−3 10−2 10−1 100 101 102

Ln (1052 erg s−1)

ne ¯ ne nx

Anatomy of the Neutrino Signal

21

Core deleptonization
Deleptonization burst
Accretion phase
Mantle contraction
Core Cooling
Neutrinosphere recession

LR and Reddy ‘15

SLIDE 22

10−1 100 101 102

tpost-bounce (s)

101 102

R (km)

Rpns Rn

(MeV)

10−1 100 101 102 10−4 10−3 10−2 10−1 100 101 102

Ln (1052 erg s−1)

ne ¯ ne nx

Anatomy of the Neutrino Signal

22

Core deleptonization
Deleptonization burst
Accretion phase
Mantle contraction
Core Cooling
Neutrinosphere recession

LR and Reddy ‘15

SLIDE 23

Early Time Neutrino Emission

23

0.2 0.4 1.2 1.4 0.8 1.0 0.6

ξ1.75

1D Study of progenitor

dependence of neutrino emission

pre-Explosion neutrino

emission driven by accretion

Progenitor core structure

determines accretion rate

Dependence on nuclear EoS

via neutron star compactness

100 200 300 400 20 40 60 80 100 120 140

Lν [1051 erg s-1]

100 200 300 400

t-tbounce [ms]

10 15 20 25 30

<E> [MeV]

10 15 20 25 30

νe νe

O’Connor & Ott (2013)

SLIDE 24

Late Time Neutrino Emission

Kelvin-Helmholtz evolution of the neutron star mediated by

neutrinos

Coupled neutron star structure and neutrino transport
Sensitive to dense matter equation of state, neutrino oscillations
Possibly cleaner problem than explosion mechanism

See e.g. Burrows & Lattimer ’86, Pons et al. ‘99, Huedepohl et al. ‘10, Fischer et al. ’10, LR ’12, Nakazato ‘13

0 10−1 100 101 102

tpost-bounce (s)

0.0 0.2 0.4 0.6 0.8 1.0

Fractional Energy

16

PNS Cooling

SLIDE 25

Simple Prescription for Explosion in 1D

Once supernova shock passes fixed mass shell, remove all
f the overlying mass and replace with a boundary

condition

Drawback: Abrupt end to accretion
Makes baryonic mass of remnant a free parameter, but we

don’t know it anyway without realistic explosion model

25

Perform an (inverse) mass cut

SLIDE 26

Long Term PNS Evolution

26

0.0

0.2 0.4 0.6 0.8 1.0 1.2 1.4

Enclosed Mass (M)

101 102 103

r (km)

0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4

Enclosed Mass (M)

105 104 103 102 101

nb (fm3)

25 ms 100 ms 1 s 10 s 70 s

ESN ⇠ 3GM2

pns

5rNS ⇡ 3⇥1053 erg ✓Mpns M ◆2 ⇣ rNS 12km ⌘1 .

SLIDE 27

Long Term PNS Evolution

27

0.0

0.2 0.4 0.6 0.8 1.0 1.2 1.4

Enclosed Mass (M)

2 4 6 8

s (kb/baryon)

0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4

Enclosed Mass (M)

0.05 0.10 0.15 0.20 0.25 0.30 0.35 0.40

YL

25 ms 100 ms 1 s 10 s 70 s

τc ⇡ 2πG2

Fc2 A

β ⌧ N0 3nb π2 ∂s ∂T

kBTc R2 ' 10 s

kBTc 30 MeV hn2/3

b i

n2/3 ✓ R 12 km ◆2

See Prakash et al. ‘97

SLIDE 28

Detection Rates

28

From Shirley Li

SLIDE 29

Progenitor Dependence

29

10−1 100 101

tpost-bounce (s)

10−2 10−1 100 101 102

Lν (1052 erg s−1) ˙ Nν (1057 s−1)

10−1 100 101

tpost-bounce (s)

0.1 0.2 0.3 0.4

Ye and YL

s40 s15 1.6 Msun PNS

SLIDE 30

Progenitor Dependence

30

− −

100

ν ν ν ν ν ν

0.45 0.50 0.55 0.0 0.5 1.0 1.5 2.0 30

Inner-core Mass (M )

Sensitivity to

Progenitor variations Electron capture rate variations

sc (kb / baryon)

sc

− −

ν ν ν ν ν ν

P r

g

e n i t

r

s ( 3 2 ) s 1 2

s

1 2 W H 7 s 1 2 W H 7 + S F H

s

2 W H 7 + S F H

s

4 W H 7 + S F H

s

1 5 W W 9 5 + S F H

s

1 5 W W 9 5 + D D 2 s 1 5 W W 9 5 + T M A

From Sullivan et al. (2015) Properties of the inner core after bounce are relatively insensitive to progenitor structure

Liebendoerfer et al. 2002

SLIDE 31

Proto-Neutron Star Convection

Region of convective instability determined by the Ledoux Criterion:

10−1 100 101 102

Time [s]

101

Radius [km]

Rνe R ¯

νe

Rνx

e d s

n

), d

e

e

d
n

SLIDE 32

Black: No Convection Red: Convection

10−1 100 101 102 10−4 10−3 10−2 10−1 100 101 102

Ln (1052 erg s−1)

ne ¯ ne nx

post-bounce

10−1 100 101 102

tpost-bounce (s)

4 6 8 10 12 14 16

en (MeV)

102 10−1 100 101 102

tpost-bounce (s)

0.0 0.5 1.0 1.5 2.0

s (kb/baryon)

−1

1 2

0.1 0.2 0.3 0.4

Ye and YL

Proto-Neutron Star Convection

SLIDE 33

Proto-Neutron Star Convection

Pressure derivatives are sensitive to the symmetry energy derivative: LR et al. (2012)

HIC Skin (Sn) PDR IAS

0.1 0.2 0.3 0.4 0.5

n (fm -3 )

50 100

IU-FSU GM3

30 35 Esym(n ) (MeV) 5 10 15 20 25 30 35 n0E’sym (MeV)

QMC

Esym (MeV)

B

Dependence on the EoS

SLIDE 34

Count Rate (s−1) Time (s) 10 10

1

10

1

10

2

10

3

Convection MF GM3 No Convection g’=0.6 GM3 Convection g’=0.6 GM3 Convection g’=0.6 IU-FSU

0.3 0.35 0.4 0.2 0.25 0.3 0.35 Counts (0.1 s −> 1 s)/ Counts (0.1 s −> )

∞

Counts (3 s −> 10 s)/ Counts (0.1 s −> )

∞

0.45

LR et al. (2012)

Proto-Neutron Star Convection

SLIDE 35

Opacity Dependence of Late Time Cooling

35

10−2 10−1 100 101 102

Total ne ¯ ne nx

18 10−2 10−1 100 101 102

˙ N 1057 s−1

10−2 10−1 100 101 102

Total ne ¯ ne nx

18 10−2 10−1 100 101 102

˙

57

−1

10−1 100 101

tpost-bounce (s)

0.5

2 1

˙ Nn (1057 s−1)

10 10

Ln (1052 erg s−1)

SLIDE 36

Impact of Nuclear Correlations on Neutrino Opacities

Correlations through the RPA:

0.1 0.2 0.3 0.4 0.5 0.6 0.7

nb (fm−3)

10 20 30 40 50

T (MeV)

DRPA

2

/DMF

2

0.1 0.2 0.3 0.4 0.5 0.6 0.7

nb (fm−3)

DRPA

3

/DMF

3

0.1 0.2 0.3 0.4 0.5 0.6 0.7

nb (fm−3)

DRPA

4

/DMF

4

Yνe = 0.05

0.5 1.0 1.5 2.0 2.5 3.0 3.5 4.0

See Horowitz ’93, Reddy et al. ’99, and Burrows & Sawyer ‘99

Neutrino Diffusion Coefficients

SLIDE 37

Impact of Screening

Lν,tot (1051 erg s−1)

Time (s) 10

−1

10 10

1

10 10

1

10

2

Base Convection RPA RPA + Convection

LR et al. (2012) see also Huedepohl et al. (2010)

SLIDE 38

Variations in the Interaction

Reddy et al. (1999)

Varying the axial interaction

LR et al. (2012)

SLIDE 39

The Neutrino Driven Wind

After successful core collapse supernova, hot

dense Protoneutron Star (PNS) is left behind

As neutrinos leave the PNS, they deposit

energy in material at the neutron stars surface

Drives an outflow from the surface of the

neutron star

Electron fraction is determined by the

neutrino interactions, some neutrons turned into protons and vice-versa

Possible site to make some interesting nuclei

that are not made during normal stellar evolution: r-process, light p nuclides, N = 50 closed shell nuclei Sr, Y, Zr

See Duncan et al. ‘86,Woosley et al. ’94, Takahashi et al. ‘94, Thompson et al. ’01, Metzger et al. ‘07 Arcones et al. ’08, LR et al. ’10, Fischer et al. ‘10, Huedepohl et al. ’10, Vlasov ’14, etc.

SLIDE 40

What Determines the νe Spectra?

“Neutrino sphere” is not well

defined, energy dependent, range of densities and temperature

Both charged and neutral current

reactions important to νe and anti-νe decoupling radii

Charged current rates introduce

asymmetry between neutrinos and antineutrinos

From Raffelt ‘01

SLIDE 41

0.1 1 10 Time (s) 10

54

10

55

10

56

LLepton (# s

1)

IU-FSU No MF 10

50

10

51

10

52

10

53

LE (erg s

1)

Neutrino emission w/ and w/o Nuclear Interactions

Self Energies

See LR ‘12 and Martinez-Pinedo et al. ‘12

N

S

e l f E n e r g i e s

Deleptonization

S e l f E n e r g i e s No Self Energies

Self energies shift average neutrino energies

0.1 1 10 Time (s) 6 9 12 15 <> (MeV)

e (U = 0) e x e (U=RMF)

B

SLIDE 42

2 4 6 8 10 Time (s) 0.45 0.5 0.55 0.6 Ye,NDW IU-FSU No MF

Neutrino emission w/ and w/o Nuclear Interactions

See LR ‘12 and Martinez-Pinedo et al. ‘12

SLIDE 43

Huedepohl et al. ’10 neutrino histories. Very little

nucleosynthesis. 7.5 Msun ejected.
Roberts. ’12 neutrino histories. Significant N = 50

closed neutron shell production.

Integrated NDW Nucleosynthesis

mass number abundance 50 100 150 200 250 10-9 10-8 10-7 10-6 10-5 10-4 10-3 10-2

M = Msun 1.2 1.4 1.6 1.8 2.0 2.2 2.4 solar r-abundance

Wanajo (2013)

SLIDE 44

Symmetry Energy Dependence

2 4 6 8 10 Time (s) 6 8 10 12 14 <> (MeV) 2 4 6 8 10 6 8 10 12 14 No MF (IU-FSU) IU-FSU GM3

e e

From Roberts et al. (2012) From Horowitz et al. (2012) Different equations of state

SLIDE 45

Convection increases deleptonization rate, increases Ye
Convection heats up PNS atmosphere, hotter neutrino spectra,

decreases Ye

…Convection

45

from Mirizzi, et al. (2015)

SLIDE 46

Conclusions

3D models of radiation hydrodynamic models of CCSNe starting to

become available, producing explosions

PNS convection significantly impacts the neutrino cooling timescale,

produces a break in the neutrino emission, sensitive to the nuclear EoS

Neutrino opacities especially important to the late time cooling

timescale

In particular, nuclear correlations can also leave a signature on the tail
f the neutrino signal
Properties of the neutrinos can also impact nucleosynthesis near the

PNS

46