Gamma-Ray Bursts: 2. Long GRBs Brian Metzger, Columbia University - - PowerPoint PPT Presentation

Gamma-Ray Bursts:

• 2. Long GRBs

Brian Metzger, Columbia University

BATSE Bursts (from Nakar 2007) short

Gamma-Ray Burst Durations

Duration long

GRB 030329 and the Supernova Connection

Exploding “Wolf-Rayet” Star

radius R~1011 cm (3 light-seconds).

Gamma-Ray Burst Galaxies

(courtesy A. Fruchter)

⇒ Long GRBs come from the

deaths of massive Stars

GRB 030329 and the Supernova Connection

Exploding “Wolf-Rayet” Star

radius R~1011 cm (3 light-seconds).

The ‘Collapsar’ Model

“GRBs are powered by accretion onto a new formed black hole”

(Woosley 1993)

Bl Black Ho Hole Model

(Woosley 93; MacFadyen & Woosley 1999)

• Energy

Energy

• Accretion Power
• Duration

Duration - Collapse Time of Star

MacFadyen & Woosley 1999 Zhang, Woosley & Heger 2004

E ~ " jM#c 2 ~ 1051 ergs " j 10$3 % & ' ( ) * M# M! % & ' ( ) *

t ff ~ 3" 32G#$ % & ' ( ) *

1/ 2

~ 100 s

Generally inefficient: L j <10"3 ˙ M c 2

MHD Powered Jets (e.g. Blandford & Znajek 1978)

Rezzolla et al. 2010

How is magnetic field generated?

credit: Stan Woosley

Mi Millisecond Ma Magnetar Mo Model

• Energy

Energy

• Rotation
• Luminosity -

Luminosity - Dipole Radiation

• Duration

Duration - Spin-Down Time

Bucciantini, Metzger et al. 2011

Erot ~ 12 I"2 ~ 3 #1052 ergs P1 ms

( )

$2

Magnetar Wind Surrounding Star

Ω

(e.g. Usov 1992; Metzger et al. 2011)

Lsd = µ2"4 c 3 # 6 $1049 P 1 ms % & ' ( ) *

+4

Bdip 1015 G % & ' ( ) *

2

erg s-1

"sd = Erot Lsd #10 P 1 ms $ % & ' ( )

2

Bdip 1015 G $ % & ' ( )

• 2

min

Magnetars: Super-Magnetized Neutron Stars

• Surface Magnetic Field

1014-1015 G (would erase your

credit card at distance of Sun).

• Observationally

classified as “soft gamma- ray repeaters” and “anomalous X-ray pulsars”

• “Giant Flares” every

~10-100 years.

• 12 in Milky Way
• Age: 103 -104 yrs
• rotation period P ~ secs

SGR1806-20 Giant γ-Ray Flare in December 2004

What Produces Magnetar Fields?

E rot = 1 2 I"2 ~ 3 #1052 P 1 ms $ % & ' ( )

*2

ergs

Ro ~ 1 for P ~ 1 ms

"E rot = B2 8# $ 4# 3 Rns

3 % Beq ~ 1017 "&

&/2 ' ( ) * + ,

2

P 1 ms ' ( ) * + ,

• 2

G

rotational energy:

Dessart et al. 2006

All neutron stars form as hot, differentially-rotating ‘proto-neutron stars’ Field amplification:

Log(Ro " P#c)

ΔΩ

L

" ~ 4#R2$Vc 3, lP ~ 0.1RNS

% c ~ lP Vc ~ 1 ms lP 0.1RNS & ' ( ) * + RNS 12 km & ' ( ) * +

5/ 3

$ 1014g cm-3 & ' ( ) * +

1/ 3

L

"

1052erg s-1 & ' ( ) * +

,1/ 3

Pizzolato et al. 2003 Rossby Number

• Shear instabilities (talk by Zrake)
• Magneto-rotational instability
• α-Ω dynamo (Thompson & Duncan 1993)

Vc

Magnetic activity of late type stars

Core Collapse with Magnetic Fields & Rotation

(e.g. LeBlanc & Wilson 1970)

Neutron Star Mass

˙ M

IN

˙ M

OUT

Time

“Failed Collapsar”

Neutrinos heat proto-NS atmosphere (e.g. νe + n ⇒ p + e-)

⇒ drives wind behind outgoing supernova shock (e.g. Qian & Woosley 96)

Neutrino Driven Wind

Burrows, Hayes, & Fryxell 1995

˙ M ~ 10"4 L

#

1052erg s-1 $ % & ' ( )

5/ 3

*# 10 MeV $ % & ' ( )

10/ 3

M! s"1

⇒ crucial to baryon loading

Neutrinos heat proto-NS atmosphere (e.g. νe + n ⇒ p + e-)

⇒ drives wind behind outgoing supernova shock (e.g. Qian & Woosley 96)

Neutrino-Heated Wind

Before SN Shock Launch After Shock Launch

Neutrino Driven Wind

Burrows, Hayes, & Fryxell 1995

˙ M ~ 10"4 L

#

1052erg s-1 $ % & ' ( )

5/ 3

*# 10 MeV $ % & ' ( )

10/ 3

M! s"1

⇒ crucial to baryon loading

Effects of Strong Magnetic Fields

“Helmet - Streamer”

Ω

• Microphysics (EOS, ν Heating & Cooling)

– Important for B ≥ 1016 G (Duan & Qian 2005)

Effects of Strong Magnetic Fields

“Helmet - Streamer”

B2 8" > 12 # vr

2

Outflow Co-Rotates with Neutron Star when

RA Rheat Ω

Top View

• Microphysics (EOS, ν Heating & Cooling)

– Important for B ≥ 1016 G (Duan & Qian 2005)

• Magneto-Centrifugal Slinging

(Weber & Davis 1967; Thompson, Chang & Quataert 2004)

⇒

Magneto-Centrifugal Acceleration (“Beads on a Wire”)

Regimes of Magnetized PNS Winds (B = 3×1014 G)

Neutrino Luminosity (1051 erg s-1)

Thermally-Driven

Magnetically-Driven, Ultra-Relativistic

Magnetically-Driven, Mildly Relativistic

Metzger, Thompson, Quataert 2007

Rotation Period (ms)

Regimes of Magnetized PNS Winds (B = 3×1014 G)

Neutrino Luminosity (1051 erg s-1)

Thermally-Driven

Magnetically-Driven, Ultra-Relativistic

Magnetically-Driven, Mildly Relativistic

Metzger, Thompson, Quataert 2007

Rotation Period (ms)

t ~ 100 s t ~ 1 s

Regimes of Magnetized PNS Winds (B = 3×1014 G)

Neutrino Luminosity (1051 erg s-1)

Thermally-Driven

Magnetically-Driven, Ultra-Relativistic

Magnetically-Driven, Mildly Relativistic

Metzger, Thompson, Quataert 2007

Rotation Period (ms)

t ~ 100 s t ~ 1 s

Evolution of Proto-Magnetar Outflows

(BDM et al. 2007, 2011) Initial rotation period P0 , dipole field Bdip & obliquity θdip

NS Cooling Luminosity

3D Magnetosphere Geometry

(e.g. Bucciantini et al. 2006; Spitkovsky 2006)

Calculate:

Wind Power ˙ E (t), Mass Loss Rate ˙ M (t), " 'Magnetization' #(t) ~ ˙ E ˙ M c 2 = $max(t)

In terms of

Roberts 2012

Neutrino Cooling Evolution

"0 ~ #

max =

˙ E ˙ M c2 $ B2%4 L&

5/3T10/3

spin - down power ˙ E

iso /1050 erg s-1

magnetization "0 ~ #

max

Example Solution increases as magnetar cools

RADIO X-RAYS OPTICAL

SNR PWN PULSAR

Multi-Wavelength Crab Nebula

3C58 (Chandra)

Collimation via Stellar Confinement

Collimation via Stellar Confinement

RADIO X-RAYS OPTICAL

SNR PWN PULSAR

Multi-Wavelength Crab Nebula

3C58 (Chandra)

Ω Ω

Supernova remnant elongated by anisotropic magnetic stresses in pulsar nebula? (Begelman & Li 1992)

Outgoing SN shock VSN ~ 0.1 c

Fast Magnetar Wind Vw ~ c Outgoing SN shock VSN ~ 0.1 c

Outgoing SN shock VSN ~ 0.1 c

Outgoing SN shock VSN ~ 0.1 c

Jet Formation via Stellar Confinement

(Bucciantini et al. 2007, 08, 09; cf. Uzdensky & MacFadyen 07; Komissarov & Barkov 08)

Zoom Out

Jet power & mass-loading match (on average) outflow from central magnetar

2D 3D

Porth, Komissarov, & Keppens 13

Kink Instability

Outflow becomes relativistic at t ~ 2 seconds; Jet breaks out of star at tbo ~ R/βc ~ 10 seconds

Jet Break-Out Non-Relativistic (σ0 < 1) Relativistic (σ0 > 1)

"0 ˙ E

iso /1050 erg s-1

Outflow becomes relativistic at t ~ 2 seconds; Jet breaks out of star at tbo ~ R/βc ~ 10 seconds

Jet Break-Out Non-Relativistic (σ0 < 1) Relativistic (σ0 > 1)

"0 ˙ E

iso /1050 erg s-1

Jet Break-Out

←GRB→

Central Engine

GRB / Flaring Relativistic Outflow (Γ >> 1) Afterglow

• 1. What is jet’s composition? (kinetic or magnetic?)
• 2. Where is dissipation occurring? (photosphere? deceleration radius?)
• 3. How is radiation generated? (synchrotron, IC, hadronic?)

~ 107 cm Photospheric IC

GRB GRB Em Emissi ssion

• n - What

hat, Where, here, How? How?

Metzger et al. 2011

˙ E

jet

Jet Break-Out

Optically-Thick Optically-Thin

Time-Averaged Light Curve

Photospheric Dissipation (IC)

Metzger et al. 2011

˙ E

jet

Jet Break-Out

Optically-Thick Optically-Thin

Time-Averaged Light Curve Hot Electrons ⇒ IC Scattering (γ-rays)

and Synchrotron (optical) E FE (1050 erg s-1) Spectral Snapshots t ~ 30 s E (keV) t ~ 15 s

Synch IC Tail BB

End of the GRB = Neutrino Transparency?

Ultra High-σ Outflows

⇒

• Acceleration is Inefficient

(e.g. Tchekhovskoy et al. 2009)

• Internal Shocks are Weak

(e.g. Kennel & Coroniti 1984)

• Reconnection is Slow

(e.g. Drenkahn & Spruit 2002)

TGRB ~ Tν thin ~ 20 - 100 s

"0

˙ E

iso /1050 erg s-1

←GRB→

baryons e-/e+ pairs

Steep Decline

End of the GRB = Neutrino Transparency?

Ultra High-σ Outflows

⇒

• Acceleration is Inefficient

(e.g. Tchekhovskoy et al. 2009)

• Internal Shocks are Weak

(e.g. Kennel & Coroniti 1984)

• Reconnection is Slow

(e.g. Drenkahn & Spruit 2002)

TGRB ~ Tν thin ~ 20 - 100 s

"0

˙ E

iso /1050 erg s-1

←GRB→

baryons e-/e+ pairs

Low plateau efficiency consistent with Lu & Zhang 2014

←GRB→

τSD

e.g. Zhang & Meszaros 2001; Troja et al. 2007; Yu et al. 2009; Lyons et al. 2010

Late-Time Spin-Down

˙ E

iso /1050 erg s-1

←GRB→

τSD

Willingale et al. 2007

`Plateau’ Time after trigger (s)

X-ray Afterglow

e.g. Zhang & Meszaros 2001; Troja et al. 2007; Yu et al. 2009; Lyons et al. 2010; Rowlinson et al. 2010, 2013; Gompertz et al. 2013

Late-Time Spin-Down

˙ E

iso /1050 erg s-1

A Diversity of Magnetar Birth P0 (ms) Bdip (G)

Classical GRB

Eγ~1050-52 ergs, τjet < 1, Γ ~ 102-103

A Diversity of Magnetar Birth P0 (ms)

Classical GRB

Eγ~1050-52 ergs, τjet < 1, Γ ~ 102-103 Thermal-Rich GRB (XRF?) Eγ~1050 ergs, τjet ~ 1, Γ < 10

L

• w

L u m i n

• s

i t y G R B

Bdip (G)

A Diversity of Magnetar Birth

Buried Jet

P0 (ms)

Classical GRB

Eγ~1050-52 ergs, τjet < 1, Γ ~ 102-103

Very Luminous SNe?

(Kasen & Bildsten 10; Woosley 10)

Thermal-Rich GRB (XRF?) Eγ~1050 ergs, τjet ~ 1, Γ < 10

L

• w

L u m i n

• s

i t y G R B

Bdip (G)

Galactic Magnetars?

Observational Tests & Constraints

• Max Energy* -

EKE+Eγ < 3×1052 ergs

*subject to uncertainties in afterglow modeling.

(e.g. Zhang & MacFadyen 09).

• Long GRB always accompanied by bright, energetic
• Consistent with observations thus far (Woosley & Bloom 2006).
• Γ increases during GRB and correlates with Eγ
• translate jet luminosity/magnetization into unique prediction for gamma-

ray light curves and spectra. Cenko et al. 2011

Courtesy A. MacFadyen

The Proto-Magnetar Model for Gamma-Ray Bursts

Brian Metzger

NASA Einstein Fellow Princeton University

In collaboration with

Eliot Quataert (UC Berkeley) Todd Thompson (Ohio State) Dimitrios Giannios (Princeton) Niccolo Bucciantini (Nordita) Jon Arons (UC Berkeley)

University of Minnesota Astronomy Colloquium, October 22, 2010

Long GRBs = Massive Stellar Death

GRB 030329 ⇔ SN 2003dh (BL Type Ic)

(Paczynski 98, Galama et al. 98, Bloom et al. 99, Pian et al. 06, Modjaz et al. 06, Woosley & Bloom 06)

HST, Fruchter+ 2006 Stanek+ 2003

Broad Spectral Lines ⇒ ESN ~ 1052 ergs (“Hyper-Novae”)

Average Duration ~10 - 30 s Typical Variability ~ 1 s

Long Duration Gamma-Ray Bursts

Light Curves Spectra

Broken power-law (‘Band’) spectrum w Epeak~300 keV

Amati-Yonetoku Relations:

Epeak ∝ Eiso

0.4,Liso 0.5

Epeak

• Energies - Eγ ~ 1049-52 ergs
• Rapid Variability (down to ms)
• But Most `Power’ on Timescales ~ 1s
• Duration - Tγ ~10-100 seconds
• Steep Decay after GRB
• Ultra-Relativistic, Collimated Outflow Γ ~ 100-1000
• Association w Energetic Core Collapse Supernovae
• Late-Time Central Engine Activity (Plateau & Flaring)

Constraints on the Central Engine

BH NS

versus

Canonical GRB Lightcurve

Nakar 07

Courtesy A. MacFadyen

The Fates of Massive Stars (Heger et al. 2003)

Assumes neutrino-powered supernova with energy ~ 1051 ergs

Th The Co Collapsar (F (Failed Supernova) Model (Woosley 93)

• Energy

Energy

• Accretion / Black Hole Spin
• Duration

Duration - Stellar Envelope In-Fall

• Hyper-Energetic

Hyper-Energetic SNe Ne - Delayed Black Hole Formation

• r Accretion Disk Winds
• Late-Time Activity -

Late-Time Activity - Fall-Back Accretion

MacFadyen & Woosley 1999 Zhang, Woosley & Heger 2004

(e.g. Aloy et al. 2000; MacFadyen et al. 2001; Proga & Begelman 2003; Takiwaki et al. 2008; Barkov & Komissarov 2008; Nagataki et

• al. 2007; Lindler et al. 2010)

What Distinguishes GRB Supernovae?

Soderberg et al. 2006, 2007, 2009

“GRB-SNe are not clearly distinguished from ordinary SNe Ibc either by optical luminosity or photospheric velocities.”

Core Collapse with Magnetic Fields & Rotation

(e.g. LeBlanc & Wilson 1970; Bisnovatyi-Kogan 1971; Akiyama et al. 2003)

Co Collapsar Re Requir irements:

• Angular Momentum

ngular Momentum

• Strong, Ordered Magnetic Field

Strong, Ordered Magnetic Field

(e.g. (e.g. Proga roga & & Begelman egelman 2003; McKinney 2006) 003; McKinney 2006)

Neutron Star Mass

˙ M

IN

˙ M

OUT

Heger et al. Black Hole-Neutron Star Dichotomy

(at Sub-Solar Metallicity)

Millisecond Magnetar Model (Usov 92; Thompson 94)

• Rapid Rotation ⇔ Efficient α-Ω Dynamo ⇔ Strong B-Field at P ~ 1 ms

(Duncan & Thompson 1992; Thompson & Duncan 1993)

E Rot " 2 #1052 P 1 ms $ % & ' ( )

*2

ergs E

"

#1049 P 1 ms $ % & ' ( )

*4

BDip 1015 G $ % & ' ( )

2

ergs s-1

Millisecond Magnetar Model (Usov 92; Thompson 94)

• Rapid Rotation ⇔ Efficient α-Ω Dynamo ⇔ Strong B-Field at P ~ 1 ms

(Duncan & Thompson 1992; Thompson & Duncan 1993)

E Rot " 2 #1052 P 1 ms $ % & ' ( )

*2

ergs E

"

#1049 P 1 ms $ % & ' ( )

*4

BDip 1015 G $ % & ' ( )

2

ergs s-1

Magnetar

Westerlund I: O7 Stars still present!

Muno +06

…and can have massive progenitors

SGR1806-20 Giant γ-Ray Flare in December 2004

Galactic Magnetars exist…

• Neutrinos Heat Proto-NS Atmosphere (e.g. νe + n ⇒ p + e-)

⇒ Drives Thermal Wind Behind SN Shock (e.g. Qian & Woosley 96)

Key Insight :

(Thompson, Chang & Quataert 04)

Neutron Stars are Born Hot, Cool via ν-Emission: ~1053 ergs in τKH ~ 10-100 s

Burrows, Hayes, & Fryxell 1995

Effects of Strong Magnetic Fields

“Helmet - Streamer”

Ω

• Microphysics (EOS, ν Heating & Cooling)

– Important for B ≥ 1016 G (Duan & Qian 2005)

Effects of Strong Magnetic Fields

“Helmet - Streamer”

B2 8" > 12 # vr

2

Outflow Co-Rotates with Neutron Star when

• Magneto-Centrifugal

Acceleration (“Bead on a Wire”)

• Enhanced Mass Loss

Rate

⇒ RA Rheat Ω

Top View

• Microphysics (EOS, ν Heating & Cooling)

– Important for B ≥ 1016 G (Duan & Qian 2005)

• Magneto-Centrifugal Outflows

(Weber & Davis 1967)

Example Thermally-Driven Wind

˙ M ~ 10"4 M! s"1 L" e ~ 8 #1051 ergs s-1; B0 =1013 G; P =100 ms (sound speed) Rs

Magnetically-Driven Wind

˙ M ~ 3"10#3 M! s#1 L" e ~ 8 #1051 ergs s-1; B0 =1015 G; P =1.2 ms RA Rs

Evolutionary Wind Models (BDM et al. 2007, 2008, 2010)

Initial Rotation Period P0 , Dipole Field Strength Bdip & Obliquity θdip

NS Cooling

3D Magnetosphere Geometry

(e.g. Bucciantini et al. 2006; Spitkovsky 2006) (Pons+99; Hudepohl+10)

Calculate:

Wind Power ˙ E (t), Mass Loss Rate ˙ M (t), " 'Magnetization' #(t) ~ ˙ E ˙ M c 2 = $max(t)

In terms of

" ~ #max = ˙ E ˙ M c2 $ B2%4 L&

5/3T10/3

˙ E

iso

" 0

RADIO X-RAYS OPTICAL

SNR PWN PULSAR

Multi-Wavelength Crab Nebula

3C58 (Chandra)

Collimation via Stellar Confinement

Collimation via Stellar Confinement

RADIO X-RAYS OPTICAL

SNR PWN PULSAR

Multi-Wavelength Crab Nebula

3C58 (Chandra)

Ω Ω

Magnetic Stresses in Pulsar Nebula Elongates SN Remnant!

(Begelman & Li 1992)

1) Outgoing SN shock (vSN ~ 0.03 c) creates a cavity

1) Outgoing SN shock (vSN ~ 0.03 c) creates a cavity 2) Magnetar wind sweeps through cavity (vW ~ c)

2) Magnetar wind sweeps through cavity (vW ~ c) 3) Termination Shock & `Magnetar Wind Nebula’ Forms 1) Outgoing SN shock (vSN ~ 0.03 c) creates a cavity

2) Magnetar wind sweeps through cavity (vW ~ c) 3) Termination Shock & `Magnetar Wind Nebula’ Forms 1) Outgoing SN shock (vSN ~ 0.03 c) creates a cavity 4) Compressed Field Increases Pressure on Axis

• Assume Successful Supernova

(35 M ZAMS Progenitor; Woosley & Heger 06)

• Inner BC from Proto-Magnetar Wind

Calcs for Bdip = 3 x1015 G and P0=1 ms

Proto-Magnetar Jet Formation

(Bucciantini et al. 2007, 2008, 2009) Average Power and Mass Loss Rate of Jet Leaving Star Match Those Set by Magnetar Wind

Jet vs. Wind Power

Wind becomes relativistic at t ~ 2 seconds; Jet breaks out of star at tbo ~ R/βc ~ 10 seconds

High Energy Emission (GRB) from t ~ 10 to ~100 s as Magnetization Increases from σ0 ~ Γ ~ 30 to ~ 103

←GRB→

Acceleration in GRB Jets

Tchekhovskoy et al. 2010

• Jet’s energy at small radii is mostly

carried by magnetic field (Poynting flux)

• Magnetic energy must be transferred

to kinetic energy to reach Γ∞ ~ σ0

• Time stationary, unconfined outflows

in ideal MHD attain Γ∞ ~ σ0

1/3 << σ0

(e.g. Goldreich & Julian 1970)

• Proposed Solutions -

1) converging (parabolic) geometry 2) time-variable outflow 3) non-ideal MHD (e.g. reconnection)

• All predict power law acceleration

Γ ∝ Rα (α<1) with max Lorentz factor Γmax ~ 102-103

Magnetar Birth - A Variety of Phenomena

Buried Jet

P0 (ms) Bdip (G)

Magnetar Birth - A Variety of Phenomena

Buried Jet

P0 (ms) Bdip (G)

Classical GRB

Eγ~1050-52 ergs, τjet < 1, Γ ~ 102-103 Thermal-Rich GRB (XRF?) Eγ~1050 ergs, τjet ~ 1, Γ < 10

L

• w

L u m i n

• s

i t y G R B