GRB (Gamma Ray Bursts) 17-05-2012 Tor Vergata. The Discovery The BATSE Era The prompt event era The Beppo-SAX Era The afterglow era The Swift Era The

GRBGRB(Gamma Ray Bursts)(Gamma Ray Bursts)

17-05-2012Tor Vergata

The Discovery

The BATSE Era

The prompt event era

The Beppo-SAX Era

The afterglow era

The Swift Era

The fast response era

GRB: OVERVIEW

1967-1973 Vela satellites:

look for X and gamma rays in order to monitor compliance

with the Geneva Limited Nuclear Test Ban Treaty of 1963

(no nuclear tests in space and atmosphere)

THE DISCOVERY

Discovered intense flashes of

Gamma-rays

Origin:

Earth?

Moon?

Sun?

Cosmic? (Klebesadel et al. 1973; Strong et al. 1974)

Compton Gamma Ray Observatory (CGRO)

THE BATSE ERA

• The second of NASA's great observatories• Operational in 1991-2000

• 4 instruments covering the 30 keV – 30 GeV energy range

BATSE observed ~ 1 GRB/day

with few degree accuracy

and rapid data dissemination,

yielding a wealth of new results

THE BATSE ERA

(The prompt event era)

Main results:

• Spatial distribution

• Time duration and distribution

• Lightcurves

• Spectral analysis

THE BATSE ERA

THE BATSE ERA – Spatial distribution

Brief (1ms-100s) intense flashes of Gamma-rays

• About 1 per day

• They are isotropically distributed in the sky

• A GRB does notrepeat.

Extragalactic origin!

THE BATSE ERA – Time duration

• GRBs duration distribution is double peaked

(e.g. Briggs et al. 2002)

Long GRBs

Short GRBs

• Short GRBs are harder than Long GRBs

(e.g. Fishman & Meegan, 1995)

Two classes of sources?

Structured, peaked lightcurves

THE BATSE ERA – Lightcurves

T ~ 1 ms

Internal shock instead of external shock scenario

Central engine works intermittently,

accelerating shells of matter to

different speeds that can collide

with each other

INTERNAL vs. EXTERNAL SHOCKS

Central engine produces a

continuous injection of matter

which slows down onto the

inter-stellar medium

Variability naturally explained Strong assumptions needed

Double power-law with a peak

at ~ 100 keV (in E2•N(E) vs. E)

THE BATSE ERA – Spectra

Spectra suggest non-thermal

processes

The matter responsible for the

emission must be optically

thick until T < 50 keV, to avoid

+ e+ + e-

... But compactness problem!!

Compactness parameter (L/R) controls the processes involving

photons:

L/R

THE COMPACTNESS PROBLEM

R can be estimated by causality arguments: R < c T

T ~ 1 ms

T is the minimum variability timescale

Even assuming the GRB in the galactic halo, L/R is “large” (GRB are too

compact), and >> 1, i.e., + e+

+ e- should suppress high energy photons

Solution: relativistic motion of the emitting region

obs = rest Tobs= Trest

/ = 102 103 = [1 – (v/c)2] -1/2

THE FIREBALL MODEL

Releasing of shells of

matter and energy in

equilibrium (fireball)

Pairs annihilate when

T < 50K (rest frame).

Thermal energy

converted into kinetic

energy: acceleration

R=1014cm: shell collisions and burst production (internal shock)

R=1016cm: interaction with the inter-stellar medium (external shock): afterglow

THE Beppo-SAX ERA

Launched on April 30, 1996 and switched-off on April 30, 2002

First and last observation of a GRB on July 20, 1996 and on April 30, 2002

The WFCs localize the burst, which

is then repointed by the narrow

field instruments.

Precise localization and follow-up with ground-based telescopes: the beginning

of the afterglow era

(The (Long) GRB afterglow era)

Main results:

• The optical afterglow

• The achromatic break

• The radio afterglow

THE Beppo-SAX ERA

GRB970228: First

optical afterglow

THE Beppo-SAX ERA – The optical afterglow

(Van Paradijs, et al., 1997)

GRB970508: First

host galaxy redshift

Reshifts in the range 0.5 4.5 z D L

E = 10511054 erg!

Break observed in

the GRB lightcurves

THE Beppo-SAX ERA – The achromatic break

Same time, all

wavelengths

Optical wavelengts X-ray wavelengts

Interpretation: GRB emission is collimated in JETS

COLLIMATION OF THE EMISSION

t1 t2 t3 t4 t5

F

tt4

Jet

Spher.

(t4)= The break must be achromaticThe break must be achromatic

decreases with time (slow down of the ejecta)

1. In case of collimated emission a break must be observed when =2. In case of spherical emission, the break must not be observed

COLLIMATION OF THE EMISSION

Frail et al. (2001)

can be computed and the intrinsic luminosities become: E = 510501051 erg

Unique energy

production

mechanism!

THE Beppo-SAX ERA – The radio afterglow

Short timescale variability at the

beginning of the observations

followed by a stable emission

GRB970508GRB970508

Interpretation: scintillation of radiation coming from

point-like sources and crossing the inter-stellar medium

Scintillations determine the size of the source in a model independent way. The size (~1017cm) is in a perfect agreement with the prediction of the Fireball model.

THE GRB PROGENITOR

NS-NS (BH-NS & BH-WD) travel far from their formation sites before producing GRB’s =>“clean environment”

Hypernovae/collapsar evolve much faster, going off in their Formation site =>“mass-richenvironment”

THE GRB PROGENITOR

NS-NS (BH-NS & BH-WD) merging ruled out for threereasons:

1. Must occur far from the star formation region

2. Need a lot of time to complete

3. Clean environments are expected

Long GRBs occur in dense environment, in the centre of blue, star-forming galaxies and up to high redshifts

What about Short GRBs? A different class of sources?

THE SWIFT ERA

Sat. 20 Nov 2004, 17:16 Sat. 20 Nov 2004, 17:16

GMTGMT

Swift lift-offSwift lift-off

BAT (Burst Alert Telescope): CZT detector 15-150 keV, detects >100 bursts per yearXRT (X-Ray Telescope): CCD detector0.2-10 keV, 5” FWHM resolutionUVOT (UltraViolet-Optical Telescope):150-10 nm, 0.3” FWHM resolution

Automated, fast slewing repointing of GRBs:T = 20 100 s, rapid broadcast of coords

(The fast response era)

Main results:

• The Short GRB afterglows

• The redshift distribution

• The “prompt” X-ray afterglow

• The rapid optical/IR follow-up

THE SWIFT ERA

z = 0.16

THE SWIFT ERA – The Short GRB afterglow

Covino et al. 2006

Short GRBs occur later than Long GRBs, in the outskirts of early-tipe

galaxies:SN/BH merging looks promising for Short GRBs

GRB050709

THE SWIFT ERA – The redshift distribution

The GRB redshift sample

has been considerably

increased and pushed to

higher redshifts

z = 6.3

Fast response to GRB events

allow us to study Star

Formation history and

absorption/dust evolution

systematically and up to high

zKawai et al. 2005

THE SWIFT ERA – The “prompt” X-ray afterglow

Prompt and afterglow emission (mainly) due to synchrotron processes, with

the contribution coming from other mechanisms (compton) to explain high

energy photons and hard spectra

Before Swift: X-ray

flux

time

WFC GRB Prompt

observation ( - hard X

band)

???

LECS/MECS 0.5-10

keV follow-up

THE SWIFT ERA – The “prompt” X-ray afterglow

In 50% of the

Swift GRBs: X-ray

flux

time

BAT GRB Prompt

observation (hard X band)

XRT 0.2-10 keV follow-up

1) High latitude emission

1

2) Late time activity ???

2

3) Late internal shocks

3

33

THE SWIFT ERA – The rapid optical/IR follow-up

In general, fast and precise identification of the GRB afterglow and fast

distribution of the coordinates have opened a wealth of new possibilities to

study the GRB phenomenon. GRB afterglows can be studied minutes later

the prompt event.

Lightcurves/photometry/spectroscopy of the continuum emission:

help to understand the physics responsible for the emission and put

constraints on the nature of the GRB progenitors.

Absorption/emission spectroscopy:

study of the properties of the host galaxy up to very high redshifts (z > 6)

• absorption due to dust along the line of sight• absorption due to gas along the line of sight

HIGH-zHIGH-z ENVIRONMENT:ENVIRONMENT:Inter-Stellar Medium of high-z galaxiesInter-Stellar Medium of high-z galaxies

TThrough absorption spectroscopyhrough absorption spectroscopy

• What can we learn from GRB/QSO optical spectra

• Why high resolution spectroscopy

• How to achieve high resolution spectra

• Basic features of the absorption spectroscopy

• Results for GRB 050730, GRB050922C, GRB 060418,

GRB080319B

OUTLINE

absorption spectroscopy is suitable to study:

1. The gas associated with the source surrounding

medium

2. The ISM of the host galaxy

3. The intergalactic matter along the line of sight

WHAT CAN WE LEARN FROM ABS. SPECTRA

GRB explosion site

Circumburstenvironment

To Earth

Host gasfar away

Intergalactic matter

1. Gas associated with the surrounding medium


Absorption lines of the gas surrounding the emitting source give information on its composition, density, temperature, velocity and distance from the central source.

Such parameters can put strong constrains on the models for the GRB progenitors and QSOs

Example: GRB021004

• Large velocity dispersion

( 3000 km/s)

• constant ionization parameter

High velocity wind from a Wolf-Rayet star progenitor instead of supernova remnant scenario Fiore, D’Elia, Lazzati et al. 2005

1. Gas associated with the surrounding medium


Example -2 : Broad Absorption Lines QSOs

probably associated with outflows from nuclear region


BAL QSOs ~15-20% of radio-quiet AGNs: evolution vs. geometry

2. The ISM of galaxies along the line of sight


The ISM gives us precious information on the metal enrichment history of the galaxies, which in turn is linked to the mass function evolution.

To now, metal enrichment in galaxies at high z has been studied using:

• Lyman Break Galaxies

• Galaxies along the line of sight of quasars (Damped Lyman- systems)

The first class cannot be representative of the true galaxy population

The second one is entangled by selection effects: the radiation from the QSOs probes preferentially the halos of the galaxies.

2. The ISM of the host galaxy along the line of sight


GRBs provide an independent way of studing the metal enrichment of galaxies at z > 1.

Advantages:

• No luminosity bias

• Probing central galaxy regions

• ISM can be studied up to higher redshift than DLA systems.

GRB host appear to be more metal rich than DLA systems (Savaglio 2005)

3. The Inter-Galactic Medium along the line of sight


QSOs and GRBs spectra can be used to probe the Ly- forest and the high-z intergalactic medium.

Before GRBs, IGM has been studied using the absorption systems along the QSO sightline only (quasar forest).

Using GRBs as remote beacons to study the IGM brings the following advantages:

• The proximity effect can be limited, since GRBs do not affect the IGM, as QSOs.

• The analysis can be pushed up to higher redshifts, where QSOs are not yet formed

WHY HIGH RESOLUTION SPECTROSCOPY

1. High resolution spectroscopy can disentangle the GRB surrounding medium in components, allowing a more accurate study.

GRB 021004 FORS1 R=1000 CIV z = 2.296 and z = 2.328

GRB 021004 UVES R=40000CIV z=2.296 e 2.328


2. High resolution spectroscopy is necessary to disentangle the ISM component from the absorption coming from the GRB surroundings.

GRB 050922C


GRB 050730

3. High resolution spectra provide precise dn/dz counts to put constrains on hierarchical clustering models.

HOW TO ACHIEVE HIGH RESOLUTION SPECTRA

GRB flux drops as t-1.

We need optical magnitudes < 19 19.5 in order to obtain high resolution spectra with good signal to noise ratio in a reasonable amount of time.

GRB020813: z=1.245 - 24 hours after the GRB; R=20.4

GRB021004: z=2.328 - 12 hours after the GRB; R=18.6


Swift satellite locates GRBs witharcsec precision in a few tens of seconds

GRB Coordinate Network (GCN) releases these positions in a few seconds

VLT Rapid Response Mode (RRM) allows to point such coordinates in about 8 minutes


UVES high resolution spectroscopyUVES high resolution spectroscopyUVES (Ultraviolet-visual echelle spectrograph) operate with high efficiency from the atmospheric cut-off at 300 nm to the long wavelength limit of the CCD detectors (about 1100 nm). The light beam from the telescope is splitted into two arms (UV to Blue, and Visual to Red). Arms can be operated separately or in parallel via a dichroic beam splitter.

Instrument mode λrange(nm) Maximum resolution(λ/Δλ) Covered λrange Magnitude limits

Blue arm 300-500 80,000 80 17-18

Red arm 420-1100 110,000 200-400 18-19

Dichroic #1 300-400 80,000 80 17-18

500-1100 110,000 200 18-19

Dichroic #2 300-500 80,000 80 17-18

600-1100 110,000 400 18-19

ABSORPTION SPECTROSCOPY

Basic principles

When a radiation beam encounters a

gas cloud along its path, wavelengths

related to the atomic and ionic levels

of the elements are absorbed.

N.B.: to observe a specific

transition, the corresponding

energy level must be

populated!

• 911.27 Å represents the

Lyman limit. Energies higher

than this limit are

continuously absorbed by H

atoms.


Basic principles

An example: the Hydrogen atom

n = 1

n = 2

Lyman

Incident radiation

n = 3

n = 4

Balmer

Paschen

n = ∞

1215.67Å

911.27 Å ……

Hydrogen atom

• Transitions from the n=1

level originate the Lyman

lines, n=2 the Balmer lines,

n=3 the Paschen lines and

so on…

GRB 050730, z=3.968

Lyman

limit

Ly-

Ly-

Ly-


The probability for an atom or ion to change its state depends on the

nature of the initial and final state wavefunctions, how strongly light can

interact with them, and on the intensity of any incident light.

To a first approximation, transition strengths are governed by

selection rules which determine whether a transition is allowed or

disallowed.

More quantitatively, transition strengths are usually described in

terms of the Einstein coefficients (A and B) or the oscillator strength

(f).

Transition strength

Bij = 83R2/3hgi fij Bij

0 < fij < 1i

j

Photoabsorption Incident photon

R = <i|u|j>


•The equivalent width (W) of a

line is the wavelength interval for

which the continuum and line

energies are equivalent.

Equivalent width and column densities

•The column density of a line is

the density projected along the

line of sight of the atoms or the

ions generating the line itself.

•This quantity can be estimated

using the Curve of Growth, which

describes how the line strength

(W) increases with the optical

depth


• The Curve of Growth strongly depends on the doppler parameter

of the line.

Line fitting and column densities

• An alternative and more accurate method to compute the column

densities is to fit the line using a Voigt profile (the spectral line shape which

results from a convolution of independent Lorentzian and Doppler line

broadening mechanisms).

We need to know the oscillator strength of the transition.

Output of the fit are the column density and

the doppler parameter.


Summing all the column densities of ions of the same atom yields the total

abundance of such an atom.

The ratio of the abundances of the metals with respect to Hydrogen is called

metallicity.

•The metallicity is the main tool to investigate the metal enrichment history of

the galaxies, which is linked to the mass function evolution.

•The relative abundances of different atoms can give information about the

dust content of the galaxies.

•The comparison between the relative abundances of different ions of the

same atom and photoionization models yields the ionization state of the gas,

and can put constraints on its origins.

What can we learn from column densities?


Fine structure features: The gross structure of an atom is due to the principal quantum number

n, giving the main electron shells of atoms. However, electron shells

exhibit fine structure, and levels are split due to spin-orbit coupling (the

energy difference

between the

electron spin

being parallel or

antiparallel to

the electron's orbital

moment).

Fine structure splitting

First fine structure excited level


Fine structure in absorption spectroscopy:

• Optical – UV incident radiation coming from a

background source collides on an intervening

cloud of gas.

• If the intervening gas is

composed by atomic species

whose ions have been

previously excited to the fine

structure levels, fine structure

lines with * > are observed.n

n + 1

Photoabsorption line ()

Fine structure

line (* )

Incident radiation

J=1/2

J=3/2


How to populate fine structure excited levels:

1. Collisional processes:

2. Radiative processes:

n

n + 1

Photoexcitation

Radiative de-excitation

Incident UV radiation

J=1/2

J=3/2

2a. Indirect UV pumping

J=9/2

J=7/2

J=5/2

J=3/2

J=1/2

2b. Direct IR pumping

Incident IR radiation

Selectionrule: J=0,±1

(Si II, C II) (Fe II)

Incominge-

(O I)J=0

J=1

J=2

n

n

Detailed balance equation for a two levels system:

n: density of the states - w: radiative terms - Q: collisional terms

Fine structure, assuming electron-ion collisions is main process:

(For C II) (For Si II)

ne: electron density - T: temperature - N: density of the states

INFORMATIONS ON T AND ne can be obtained.

If indirect UV pumping is instead at work, we can gather

informations on the strength of the radiation field G distance


Why studying fine structure absorption features

• Absorbing systems (host + intervening)

• Host gas separation in components

• Fine structure absorbing features

• Constraining physical parameters of the gas with fine structure

• Distance of the gas from the GRB explosion site

• Metallicity

GRB 050730: ANALYSIS AND RESULTS

3000 s Dichroic 1 3000 s Dichroic 24 hr after the GRB

LIGHT CURVE AND UVES/VLT OBSERVATIONS

Five intervening absorbers identified:

GRB 050730: ABSORBING SYSTEMS AND LINE FITTING

z1 = 3.967 (GRB host)

z2 = 3.564

z3 = 2.262

z4 = 2.253 (d)

z5 = 1.772 (d)

The main system presents 5 ( +1) components

GRB 050730: ABSORBING SYSTEMS AND LINE FITTING

C IV: 5 components1) +32.6 2) +2.4 3) -44.0 4) -90.25) -154.6

Si IV: 4 components2) +2.4 3a) -44.0 a3b) -44.0 b 4) -90.2

CIV and Si IV components used as reference to fit the other ions

FINE STRUCTURE ABSORPTION FEATURES

Fine structure transitions - the ion C II

C II 1036 and C II 1335 doublets

Only components 2 and 3 are present in the excited fine

structure features


Fine structure transitions - the ion Si II and the atom O I

Si II 1260, Si II 1304, Si II 1526 doublets and O I 1302 triplet

Only components 2 and 3 are present in the excited fine

structure features

SiII*

OI*


Fine structure transitions - the Fe II

Fe II 1608 - 1611 multiplet

Only component 2 is present in the fine structure multiplet

Second component, assuming electron-ion collisions

CONSTRAINING THE PHYSICAL PARAMETERS

From C II and Si II fine structure doublets:

103<T<104 K ne > 300 cm-3

From Fe II fine structure multiplet:

T = 2600 , ne 104 106 cm-3 +3000-900

Assuming indirect UV pumping, G/G0 = 105 106 where G0 = 1.6 X 10-3 erg cm-2 s-1.

Third component, assuming electron-ion collisions

CONSTRAINING THE PHYSICAL PARAMETERS

From C II fine structure doublet:

103 < T< 104 K10 < ne < 60 cm-3

For Si II: third component is uncertain

Assuming indirect UV pumping, G/G0 = 105 (O I) and G/G0 = 106 (Si II)Third component of Si II is uncertain, but if UV pumping is at work, we shoud observe similar columns for O I* and Fe II* (not observed).

T = 10 3 K

Relative distance of the shells corresponding to the different

components

GAS DISTANCE FROM THE GRB SITE

• Component 1 is present only with very high ionization states (CIV and OVI): it experience strong radiation field and it is probably the closest to the GRB site

• Scaling arguments both in case of electron-ion collisions (d n-1/2) and indirect UV pumping (d G/G0

-1/2) suggest that the second component is closer to the GRB site by a factor from a few to a few tens with respect to the third.

A more accurate estimate of the distance of the shells from the GRB site needs the comparison of the data with a time dependent photoionization model

THE C/Fe Ratio

Average [C/Fe] ratio is 0.080.24, consistent with values predicted for a galaxy younger than 1 Gyr undergoing star formation

[C/Fe] of component 3 is 0.530.23, larger than in 2 (-0.150.13), with [C] roughly constant.

Since Fe dust grains are more efficiently destroyed than C dust grains by the GRB UV flux and blast wave (Perna, Lazzati & Fiore 2003), this suggests that component 2 is closer to the GRB than component 3.

METALLICITY

Fitting the H absorption features, the metallicity can be estimated

We used the Ly- and Ly- absorption features to constrain the H column density. We find NH = 22.050.29.

The Hydrogen line profiles are too broad to disentangle the contributions from the five components.

Metallicity values of Z 10-3 10-2 with respect to solar are obtained.

Z can be underestimated, because: • Most Hydrogen may lie in the outer regions of the host.• Heavy elements may form dust.

Heavy elements form dust grains

THE DUST DEPLETION PROBLEM

The ISM of the Milky Way shows that Zn tends to stay in the gas phase (max 20% in dust).

On the other hand, Fe tends to form dust (up to more than 99% of total).

The ratio Zn/Fe is an excellent dust depletion indicator because:

1. Fe and Zn are extreme in their refractory properties.

2. They are easy to detect.

3. They have similar formation timescales.

(Savaglio 2005)

THE DUST DEPLETION PROBLEM

Dust depletion bias can be avoided using Cr and Zn as metallicity indicators, since they do not form dust.

Z(Cr) = -1.8 0.2

Z(Zn) = -1.3 0.2

GRB060418GRB050730GRB050922C

GRB 060418

INCREASING THE GRB SAMPLE

Fine structure: UV pumping or collisions?

In GRB 060418 fine structure lines are produced by UV pumping at r 0.5 ± 0.1 kpc (Vreeswijk et al. 2007)

ne = 109 cm-3


Fine structure variability: high vs. low resolution

Vreeswijk et al. 2007


Nice but… more statistics needed!

Searching for fine structure lines disentangled in Low-Res Dessauges-Zavadsky

et al. 2006

X-SHOOTER

High efficiency spectrograph at the UT2 Cassegrain focus

Intermediate resolution (R = 4000-14000)

Wide spectral coverage (3000-25000Å)

Three arms splitting: UVB, VIS, NIR

First light: Nov 2008

In operation: Oct 2009

Suitable to:• spot GRBs up to z ~ 20• study host metallicity in a wide redshift range• follow line variations with higher temporal resolution and longer times• collect good quality GRB spectra up to R ~ 21.5-22 (short GRBs)

THE SPECTACULAR CASE OF GRB080319B

19 March 2008, 06:12:49 UT: the brightest GRB ever

• Observed before, during and after the GRB worldwide• R=5 at about 20 s and H=4.2 at about 50 s from the GRB: naked eye GRB!

UVES observations began just 8m30s after the GRB (fastest response and higest S/N ever)Two RRM and one ToO observations of the event (8m, 2h and 3h time delay)

Fine structure lines nearly disappear

in less than 2 hours (less than 1h rest frame)!

Z=0.937


Six components clearly identified: I si the closest one


Fine structure of component III and IV drops faster than that of component I

Explanations:

Component I experience higher fluxes

for longer times, i.e., is closer to the GRB

Distance of the gas from the GRB:

dI = 0.6 kpcdIII = 1.7 kpc

TIME DEPENDENT MODELING

Balance equation:

were:

h

Absorption

up

low

up up

low low

h

Stimulated emissionSpontaneous emission

h

GRB080319B - TIME DEPENDENT MODELING

Component I Component III

Documents

GRB (Gamma Ray Bursts) 17-05-2012 Tor Vergata. The Discovery The BATSE Era The prompt event era The Beppo-SAX Era The afterglow era The Swift Era The