The Formation of Neutron Stars (and Black Holes) in Binaries · The Formation of Neutron Stars (and Black Holes) in Binaries Philipp Podsiadlowski (Oxford) • the majority of massive

The Formation of Neutron Stars (and Black Holes)

in Binaries

Philipp Podsiadlowski (Oxford)

• the majority of massive stars are in interacting binaries

• the final structure and fate of massive stars is very different in

binary systems

I. Binary Interactions

II. The Fates of Stars in Binaries (vs. Single Stars)

III. Supernova Kicks

IV. Black-Hole Formation

Binary Interactions

• most stars are members of binary

systems

• a large fraction are members of

interacting binaries (30− 50%)

Sana et al. (2012):

70% for O stars with M ∼> 15M⊙

• note: mass transfer is more likely for

post-MS systems

• mass-ratio distribution:

⊲ for massive stars: masses correlated

⊲ for low-mass stars: less certain

• binary interactions

⊲ common-envelope (CE) evolution

⊲ stable Roche-lobe overflow

⊲ binary mergers

⊲ wind Roche-lobe overflow

R/ R .

radius evolution

O

Classification of Roche-lobe overflow phases

M = 5 M

2M / M = 2

1

O.1

45 %

45 %

10 %

helium ignition

carbon ignition P = 4300 d

P = 0.65 d

P = 1.5 d

P = 87 d

Case C

main sequence

Case A

(Paczynski)

100

10

1000

10 (10 yr)750

Case B

Stable Mass Transfer

• mass transfer is ‘largely’

conservative, except at very

mass-transfer rates

• mass loss + mass accretion

• the mass loser tends to lose most of

its envelope → formation of helium

stars → hydrogen-deficient

supernovae (IIb, Ib, Ic)

• the accretor tends to be rejuvenated

(i.e. behaves like a more massive star

with the evolutionary clock reset)

• orbit generally widens

Unstable Mass Transfer

• dynamical mass transfer →

common-envelope and spiral-in phase

(mass loser is usually a red giant)

⊲ mass donor (primary) engulfs

secondary

⊲ spiral-in of the core of the primary

and the secondary immersed in a

common envelope

• if envelope ejected → very close binary

(compact core + secondary)

• otherwise: complete merger of the

binary components → formation of a

single, rapidly rotating star

The Formation of NS-NS (NS-BH) Binaries

X X

XX

v∆

v∆

X

X

X

X

HeNS

Core Collapse Supernovae

Iron core collapse

• inert iron core (> MCh)

collapses

⊲ presently favoured model:

delayed neutrino heating

to drive explosion

νν ν

ν ν νν

νν

ννννννν

νν

Kifonidis

NS

Iron Core

Collapse

Electron-capture supernovae

• occurs in degenerate ONeMg core

⊲ at a critical density

(4.5× 109 g cm−3), corresponding

to a critical ONeMg core mass

(1.370± 0.005M⊙), electron

captures onto 24Mg removes

electrons (pressure support!)

→ triggers collapse to form a low-mass

neutron star

note: essentially the whole core

collapses

→ easier to eject envelope/produce

supernova

→ no significanct ejection of heavy

elements

��

��

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−−> electron−capture supernova

without dredge−up−−> larger CO core mass

in ONeMg coreo lower explosion energyo lower supernova kickso NS mass: 1.25 Msun

Second dredge−up in AGB stars (around 10 Msun)

with H envelope without H envelope

AGB envelope

dredge−up of the He core−−> lower CO core masses−−> ONeMg WD

CO core

(Podsiadlowski et al. 2004)

Binary Evolution Effects

• dredge-up in AGB phase may prevent

ONeMg core from reaching Mcrit →

ONeMg WD instead of collapse

• can be avoided if H envelope is removed

by binary mass transfer

→ dichotomous kick scenario

(Podsiadlowski, Langer, et al. 2004)

⊲ e-capture SN in close binaries → low

kick

⊲ iron core collapse → high kick

Subsequent Work: Single Stars

Arend Jan Poelarends (PhD Thesis):

• examined conditions for e-capture SNe on metallicity, wind mass loss, dredge-up

efficiency in AGB stars

• best model: no e-capture SN at solar Z

Recent Work: Binary Stars → Thomas Tauris

��

��

��

��

helium burningconvective core(growing)

No H−burning shell

shrinkingHe−burning core

without H envelope

H−burning shell

with H envelope

−−> larger CO cores with lowerC/O ratio −−> no convective carbon burninghigher entropy (more massive) iron cores

−−−> BLACK HOLE

smaller CO cores with higher

lower entropy (mass) iron cores

He−core−burning stars (M > 20 − 25 Msun)

(Brown, Lee, Heger)−−−> NEUTRON STAR (60/70 Msun?)

C/O ratio −−> convective carbon burning

Brown, Heger, Langer et al. (2001)

Carbon Burning and Final Fe CoreMasses

(Brown et al. 2001)

• late He-core burning: 12C + α becomes

dominant and determines the final 12C

fraction

⊲ stars with H-burning shell: injection of

fresh He → long 12C + α phase → low final

C fraction

⊲ stars without H-burning shell: short12C + α phase → higher final C fraction

• C-core burning:

⊲ high C fraction → convective C burning

→ higher neutrino losses → lower-entropy

cores → lower-mass O and ultimately Fe

cores → neutron stars

⊲ low C fraction → radiative C burning →

lower neutrino losses → higher-entropy

cores, etc. → black holes

Petermann, Langer & Podsiadlowski (2015)

Petermann, Langer & Podsiadlowski (2015)

LBV Supernovae from MassiveBinary Mergers

Justham, Podsiadlowski & Vink (2014)

• large number of O-star binary mergers

(Sana et al. [2012]: 20–30%)

• for sufficiently small core mass fraction

⊲ He burning in blue-supergiant phase

⊲ with relatively low-mass loss rate

⊲ transition to the red only after

He-core burning

→ possibility of SN explosion in LBV

phase

(with various amounts of H envelope

masses)

Justham et al. (2012)

• even relatively massive stars may

produce neutron stars rather than

black holes (low entropy, plus core

erosion)


Neutron Star Birth Kicks

• single radio pulsars have large space velocities (Lyne &

Lorimer; Hobbs et al. 2005): σv = 265kms−1 without evi-

dence for a low-velocity component

Evidence for Low Supernova Birth Kicks

• neutron star retention in globular clusters (e.g. Pfahl,

Ivanova)

• the existence of wide Be/X-ray binaries with low eccen-

tricities (e.g. X Per) (Pfahl)

• DNSs with low eccentricities (van den Heuvel, Dewi)

• the spin period – eccentricity relation of DNSs (Dewi)

• preference for low-kick NSs in binaries?

The origin of supernova kicks

• dramatic recent progress in neutrino-driven

core-collapse simulations

• supernova kicks produced by standing accretion shock

instability (SASI) (Blondin, Mezzacappa, Foglizzo,

Janka)

• driven by advective-acoustic instability

• l = 1 instability in two flavours:

⊲ sloshing instability (m = 0)

⊲ spiral mode (m = ±1)

• can produce kicks of a few 100 km s−1 if the collapse

phase lasts ∼> 500ms (many growth timescale)

• prediction (Podsiadlowski, Langer, et al. 2004):

⊲ large kicks for slow explosions (standard Fe cores),

⊲ low kicks for fast explosions (e-capture SN, low-mass

Fe cores)

Sloshing Instability

(l = 1, m = 0)

(Janka, Scheck, Foglizzo)

Iwakami et al. (2008)

Spiral Mode

(l = 1, m = ±1)

(Blondin, Mezzacappa)

Podsiadlowski, Langer et al. (2004)

The Bimodal Spin Distribution of NSs in Be X-Ray BinariesKnigge, Coe & Podsiadlowski (2011)

Knigge, Coe & Podsiadlowski (2011)

• spin period may be a better proxy

for NS formation channel (?)

• comparable numbers of Fe core

collapse and e-capture NSs

• Puzzle: Why is the equilibrium spin

period different for the different NS

formation channels?

⊲ different B fields? (fast vs. slow

explosion phase)

⊲ different (accretion) geometry?

(e.g. misalignment of spin)

• Be X-ray binaries may be useful for

constraining NS formation and the

formation of double NS binaries

Black-Hole Formation

• single stars (case C binaries):

⊲ fallback black holes (up to 40M⊙?;

Fryer): faint supernova, NS-type kick

⊲ ‘prompt’ black holes: no supernova,

no kick

• binaries:

⊲ may need case C (i.e. very late) mass

transfer

⊲ uncertain range of case C for massive

stars (radius evolution, wind mass

transfer?)

• accretion-induced collapse of

accreting NS in I/LMXB

⊲ for efficient accretion efficiency →

form low-mass black holes in

binaries

⊲ not observed?

• accretion-induced collapse during

in-spiral in massive envelope (e.g.

Chevalier, Fryer, Brown)

• Demorest et al. (2010): PSR

1614-2230

⊲ MNS = 1.97± 0.04M⊙,

MWD = 0.5M⊙

⊲ massive WD requires

intermediate-mass

progenitor (Li et al. 2011;

Tauris et al. 2011)

→ relatively massive NS at

birth (> 1.6M⊙)

Li, Rappaport, Podsiadlowski (2011)

• spiraling-in NS in massive envelope,

collapse to black hole if in

neutrino-dominated regime

(Chevalier, Fryer, Brown,)

• form NS+BH binaries instead of

NS+NS binary

→ implications for direct

gravitational-wave detections with

aLIGO (direct test!)

• perhaps not!

⊲ recent study by MacLeod &

Ramirez-Ruiz of Hoyle-Lyttleton

accretion with large density

gradients

⊲ streams deflected due to angular

momentum → low accretion

efficiency (∼< 0.1M⊙) → NS

survival likely

MacLeod & Ramirez-Ruiz (2015)

.Initial binary: M

1

= 14M

�

,

M

2

= 9M

�

, P

orb

= 190 d

Stable non- onservative Case

B mass transfer leaving a

helium star with M

A

He

= 4M

�

and M

0

2

= 11M

�

, P

orb

= 350 d

After �rst supernova (with

ki k v

ki k

= 50 kms

�1

):

M

0

A

= 1:337M

�

, M

0

2

= 11M

�

,

P

orb

= 8:8 yr, e = 0:82,

�v

A

sys

= 13 kms

�1

High-mass X-ray binary phase

leading to unstable mass

transfer and a

ommon-envelope and

spiral-in phase and leaving

M

0

A

= 1:337M

�

,

M

B

He

= 2:4M

�

, P

orb

= 2:8 hr

Helium star mass transfer

phase (+ spin-up of neutron

star) leaving M

A

= 1:338M

�

,

M

He

= 1:559M

�

, P

orb

= 2:6 hr

Immediately after se ond

supernova: M

A

= 1:338M

�

,

M

B

= 1:249M

�

, P

orb

= 3:3 hr,

e = 0:12, �v

B

sys

= 35 kms

�1

`Standard' Channel

X X

XX

v∆

v∆

X

X

X

X

HeNS

Double-Core Channel

v∆

v∆

X

X

X

X

He CO

Initial binary: M

1

= 11:5M

�

,

M

2

= 11M

�

, P

orb

= 3:1 yr

Unstable Case C mass

transfer: se ondary expands

to �ll its Ro he lobe

Double- ore ommon-envelope

and spiral-in phase leaving a

CO star with M

CO

= 3:0M

�

and a He star with

M

He

= 2:4M

�

, P

orb

= 3:8 hr

After �rst supernova (with

ki k v

ki k

= 300 kms

�1

):

M

0

A

= 1:337M

�

,

M

0

He

= 2:4M

�

, P

orb

= 3:3 hr,

e = 0:33, �v

A

sys

= 230 kms

�1

Helium star mass transfer

phase (+ spin-up of neutron

star) leaving M

A

= 1:338M

�

,

M

He

= 1:559M

�

, P

orb

= 2:6 hr

Immediately after se ond

supernova: M

A

= 1:338M

�

,

M

B

= 1:249M

�

, P

orb

= 3:3 hr,

e = 0:12, �v

B

sys

= 35 kms

�1


Importance of Late MassTransfer (Case C)

⊲ core evolution fixed → BH/NS

formation

⊲ late spin-up, no spin-down →

GRB progenitors

but: narrow range of predicted periods

(∼ 1%)

• at low metallicity ∼ 10%

Documents

The Formation of Neutron Stars (and Black Holes) in Binaries · The Formation of Neutron Stars (and Black Holes) in Binaries Philipp Podsiadlowski (Oxford) • the majority of massive