Univerza v Ljubljani,Fakulteta za matematiko in fiziko

Fakulteta za matematiko in fiziko

Seminar - Fourth year

Light extinction in GRB host galaxies

Author:Jure Japelj

Mentor:doc. dr. Andreja Gomboc

Abstract

In my seminar, I present two research fields of astronomy and astrophysics, namely inter-stellar extinction and γ-ray bursts, one of the most energetic explosions in the Universe, takingplace at cosmological distances. While the former is a long-time known phenomenon, its effectsbeing known for a century and a half, the latter is a relatively new branch in astrophysics,its success noticed especially in the latest decade. First, I introduce gamma-ray bursts, theirevolution and the properties of their spectra. Then I explain the basic nature of extinction, i.e.dust and gas particles, and define three basic light extinction models, based on three galaxies(Milky Way, Large and Small Magellanic Cloud). I present a few effects that gamma-ray bursthave on the surrounding dust, i.e. dust destruction. I explain how to obtain the information oflight extinction and hydrogen column density from the spectral energy distribution of the γ-rayburst afterglow. I conclude the seminar with common examples, statistical values and a specialexample of the highest detected redshift γ-ray burst.

May 8, 2010

Contents

1 Introduction 2

2 Gamma-ray bursts 22.1 The afterglow . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4

3 Light extinction 53.1 Contributors to interstellar extinction . . . . . . . . . . . . . . . . . . . . . . . . . . . 7

3.1.1 Dust . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 73.1.2 Hydrogen . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 7

3.2 Milky Way, LMC nad SMC extinction curve . . . . . . . . . . . . . . . . . . . . . . . 83.3 Dust destruction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 11

4 Extinction in GRB host galaxies 114.1 The highest redshift GRB . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 14

5 Conclusion 15

1

1 Introduction

The space between stars is not empty, but is filled with material of different origin and chemicalproperties, called interstellar medium (ISM) [1]. The ISM is of profound importance in describingstructure, dynamics and evolution of our Milky Way Galaxy, as well as other galaxies troughout theuniverse. It has a deep impact on the observations of stars and other objects that are in our Galaxy.The result is seen as an obscuration of light coming from observed objects and is called interstellarextinction. Basically it is due to summative effects of scattering and absorption of starlight. Whilewe can make direct measurements of the extinction in our own Milky Way and for a few others inour neighbourhood, the galaxies from further away cannot be examined in a traditional way. One ofsuccessful techniques is a measurement od extinction with a help of gamma-ray bursts, cosmologicalevents that happen in galaxies far away.

2 Gamma-ray bursts

Gamma-ray bursts (GRB) are the most concentrated and brightest electromagnetic explosions knownin the Universe [2]. Those are brief events, occurring at an average rate of a few per day throughoutthe universe, which for a brief period of seconds completely flood with their radiation an otherwisealmost dark gamma-ray sky.

Here I give only a brief introduction to the gamma-ray bursts. More detailed study can be foundin the seminar of Drejc Kopac (Gamma-ray bursts and their afterglows, 2009).

Although GRBs were first discovered in 1967 by the Vela satellites, they were not publicly an-nounced until 1973 [2]. These spacecrafts were flown by the U.S. Department of Defense to monitor fornuclear explosions which might violate the Nuclear Test Ban Treaty. Soon it was realized that thosegamma-ray flashes were not coming from Earth’s direction but were actually connected with a newcosmic phenomenon. The first significant step in understanding GRBs started with the 1991 launchof the Compton Gamma-Ray Observatory. The all-sky survey showed that bursts were isotropicallydistributed on the sky (Figure 1).

Figure 1: Isotropical distribution of GRB sources. Long duration, bright bursts are shown in red, whileshorter, weak bursts appear in purple. The bursts in grey had incomplete data, so the fluence could not becalculated for those. [3]

2

Without knowing the distance to the GRB source, we must give its energy output in terms ofthe total energy received per unit area of detector surface during burst (the energy flux integratedover the duration of the burst)[1]. This quantity is called fluence (S). When the distances to GRBsources have not been known yet, there were a lot of speculation whether the GRB sources areoriginating in our Galaxy† or outside it. The two scenarios give completely different results, whichthe following example will illustrate. Suppose that the fluence of a particular GRB was determinedto be 10−7J m−2. Assuming the distance to the source was to be 50,000 AU (in the Oort cloud) andthe emission of the energy from the source was isotropic, the energy of the burst would have been

E = 4πr2S = 7× 1025J.

On the other hand, if the source were located in a distant galaxy 1 Gpc away, the energy of the burstwould have been

E = 4πr2S = 1× 1045J.

The difference between the two estimates is evidently large.In 1997 satellite Beppo-SAX detected fading X-ray afterglow following GRB. Later, the discovery

of afterglow at longer wavelengths (UV-optical-IR) followed. Now measurments of redshift distancesto GRB sources were possible. Redshift of a distant source is expressed with z through equation [1]

λobsλrest

= z + 1, (2.1)

where λobs and λrest are the wavelengts in the observed and rest frame, respectively. With redshiftsmeasured, GRB sources were confirmed to be events at cosmological distances [2]. Significant ad-vances in the field were made with Swift multi-wavelength afterglow satellite, which achived the goalof accurately localizing afterglows starting less than a minute or so after the burst detection, atgamma-ray, X-ray and optical wavelength.

Figure 2: GRB duration distribution. T90 refers to time interval in which 90% of the energy is observed. [3]

After the study of thousands of events, it is clear that there are two basic classes of GRB bursts.Those that last shorter than 2 seconds are short and those that last longer are called long events. Theduration distribution of events detected by Compton Gamma-Ray Observatory is given in Figure 2.

†Beacuse of isotropical distribution of GRB sources, they would have to be very close to Earth in our Solar system,in Oort cloud for example, or nearby stars.

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The distribution indicates two fundamentally different types of gamma-ray bursts, which areassociated with two different phenomena [2] [3] [1]. For long bursts, the energy release is almostcertainly associated with the late stages of the evolution of a massive star, namely the collapse ofits core, which at least in some cases is connected with a detectable supernova. Short bursts arebelived to be associated with compact binary mergers (two neutron stars or neutron star - blackhole systems). In both cases, the central compact object is likely to be a black hole of several solarmasses. In any case, such an intense and localized release of energy implies the formation of afireball, composed of photons, e± pairs and baryons [2]. The fireball expansion has to be treated asrelativistic, in order to explain non-thermal spectra and low optical depth for high energy photons.

Figure 3: The schematic representation of the fireball evolution. Some characteristic distance and temporalvalues are also included. [3]

The evolution of the fireball is schematically shown in Figure 3. As the fireball expands, thephoton temperature inside the shell decreases, resulting in increasing fraction of annihilated pairse±. After the flow becomes optically thin for high energy photons, they can escape. To achivenon-thermal spectrum in an energetically efficient manner, the kinetic energy of the flow has to bere-converted into random energy via shocks, which are collisionless, i.e. mediated by chaotic electricand magnetic fields. The promped γ emission may be produced in two different circumstances.In the external shock model, the emission is a result of the interaction of the blast wave withthe external medium (ISM or pre-ejected stellar wind from the progenitor) and it’s properties aretherefore sensitive to the details of the medium. In the internal shock model, the energy andmomentum of the source are not constant in time. Different sublayers of the shell may collide witheachother, resulting in internal shocks before the outmost shell reaches external medium. Theseshocks can be responsible for acceleration of charged particles in inhomogenous magnetic fields, thusproducing non-thermal radiation via the synchrotron radiation and inverse Compton processes. Sincethe variations in GRB are of much shorter time scale than the whole event, it is believed, that internalcollisions are responsible for the γ-ray emmision and the external collisions for the afterglow.

2.1 The afterglow

At this point, we are mostly interested in the GRB afterglow. The afterglow spectrum at a giveninstant of time depends on the flux observed at different frequencies from electrons with comoving

4

(a) (b)

Figure 4: Syncrotron spectrum of a relativistic shock with a power-law electron distribution. (a) Fastcooling spectrum, expected at early times (initial burst). (b) Slow cooling spectrum, expected at later times(afterglow). In both figures, characteristic slopes and time evolution estimates are given. [2]

energy γemec2 and bulk Lorentz factor Γ, whose observed peak frequency is [2]

ν = ΓΓ2e

eB

2πmec. (2.2)

The deceleration of the moving shell, when it reaches external medium, is responsible for decrease inthe values of Γ and Γe, resulting in the afterglow of different frequencies. The spectrum is built ofdifferent powerlaw segments, separated by several breaks at critical frequencies of three characteristicelectron energies[2]:

• The minimum energy of the electron distribution - the injection frequency νm.

• The energy, where synchrotron cooling becomes important - cooling frequency νc.

• The frequency of synchrotron self-absorption νa.

Two different spectra exist. The first one is called the slow cooling case (νm < νc), while the secondone represents the fast cooling case (νm > νc). Both cases are shown in Figure 5. Slow coolingcase is connected with the afterglow phase. Of the three breaking frequencies, the most importantis the cooling frequency. Cooling break can be found in many detected afterglows, while the wholefrequency range, that could reveal us additional breaks, is rarely observed. Also, one has to be awareof spectral evolution, since break points will gradually move toward lower frequencies.

3 Light extinction

The Greek astronomer Hipparchus invented a numerical scale to describe how bright a star appearedon the sky[1]. He assigned an apparent magnitude m = 1 to the brightest star in the sky, and he gavethe dimmest star visible to the naked eye an apparent magnitude of m = 6†. By the modern definition,a difference of 5 magnitudes corresponds exactly to a factor of 100 in brightness. Brightness itself ismeasured in terms of radiant flux F received from the star. An observer at a distance r from a starwith luminosity L would measure a radiant flux

F =L

4πr2. (3.1)

†A smaller apparent magnitude means a brighter-appearing star.

5

Using inverse square law, we can assign absolute magnitude M to a star. That is defined to be theapparent magnitude a star would have if it were located at a distance 10pc[1]. Previously mentioneddifference of five magnitudes gives us the ratio of fluxes for two stars:

F2

F1

= 100(m1−m2)/5. (3.2)

That, combined with the definition of absolute magnitude, gives us the relation between the apparentand absolute magnitude

m−M = 5 log10

(d

10pc

). (3.3)

That equation must be modified in order to include extinction. In a given wavelength band, centeredon λ, we can write

mλ = Mλ + 5 log10

(d

10pc

)+ Aλ, (3.4)

where Aλ > 0 represents the number of magnitudes of interstellar extinction present along the lineof sight.

A few quantaties have to be explained now. First, we will introduce opacity or absorptioncoefficient κλ. The opacity is the cross section for absorbing photons of wavelength λ per unitmass of stellar material. In general, the opacity of gas is a function of its composition, density andtemperature. For a special case of a uniform gas of constant density and opacity one can write [1]

Iλ = Iλ,0e−κλρs, (3.5)

where Iλ,0 is the initial intesity of a light beam and s is a distance a beam travels through a gas cloudof density ρ and opacity κλ. The intensity declines exponentially falling by a factor of e−1 over acharacteristic distance of l = 1/κλρ. For scattered photons, the characteristic distance is in fact themean free path of the photons. Latter can be expressed differently. Let n represents a number ofparticles per cubic meter and σλ the particle cross section on which photon may be either absorbedor scattered. Then the mean free path can be written as [1]

l =1

σλn=

1

κλρ. (3.6)

Another quantity, optical depth (τλ), can be defined as

dτλ = −κλρds, (3.7)

where s is the distance measured along the photon’s path in its direction of motion. When there isonly absorption present in the gas (but no emission), it can be shown [1], that the intensity of thelight that is passing through the cloud is decreasing exponentially

Iλ = Iλ,0e−τλ , (3.8)

where optical depth τλ equals

τλ =

∫ s

0

κλρds. (3.9)

Combining (3.8) with (3.2), we can relate optical depth to the change in apparent magnitude dueto extinction

Aλ = mλ −mλ,0 = −2.5 log10

(e−τλ

)= 1.086τλ. (3.10)

The change in magnitude due to extinction is approximately equal to the optical depth along theline of sight.

6

From previous discussion, optical depth can be written as

τλ =

∫ s

0

nd(s′)σλds

′,

or, if σλ is constant along the line of sight,

τλ = σλ

∫ s

0

nd(s′)ds′ = σλNd, (3.11)

where Nd, the dust grain column density, is the number of scattering dust particles in a cylinder witha cross section of 1 m2 stretching from the observer to the source of light.

Column density and light extinction are the two quantaties usually measured which characterizethe interstellar medium. The method of measurement with the help of GRBs is explained later(chapter 4).

3.1 Contributors to interstellar extinction

3.1.1 Dust

The extinction curves and its features (especially the presence of certain dust particles) are discussedin the next section. Apparently the dust in ISM is composed of both graphite and silicate grainsranging in size from several microns down to fraction of nanometer [1] [4]. The linear polarization ofthe observed scattered light and preferable orientation of the radiation’s electric field vectors implythe non-spherical shape and special alignment of the grains. The latter is a result of the galacticmagnetic field. The grains rotate with their long axes perpendicular to the magnetic field and thelight, polarized in the plane of the rotating grain is blocked. However, the light polarized along thedirection of the magnetic field perpendicular to the line of sight is transmitted. The polarization itselfis wavelength dependant (Figure 5a) and it has been observed, that λmax at which the maximuminterstellar polarization occurs, is correlated with color excess†, and thus proportional to the averagesize of dust grains[5].

3.1.2 Hydrogen

The dominant component of ISM is hydrogen gas† in its various forms: neutral hydrogen H1, ionizedhydrogen H2 and molecular hydrogen H2. Most hydrogen is in the form of H1 in the ground state,so it is generally incapable of producing emission lines. It is also difficult to observe H1 absorptionlines, unless there are strong UV sources lying behind those clouds. However, neutral hydrogen canbe detected via unique radiation of a wavelength of 21 cm. As we know, the proton and electronmay have their spin oriented in the same or opposite direction. The hydrogen atom with the formeralignment has slightly more energy that the one with the spins being anti-parallel. A hyperfinesplitting of the orbital energy states is shown in Figure 5b right. The splitting of the ground 2Sstate is the source of the 21 cm hydrogen line. The emission of that wavelength is extremely rare.When an atom is in excited place, it can take several million years on average before that atom willemit photon and is considered forbidden by laboratory standards. However, the large number ofhydrogen atoms in ISM statistically leads to the emmision of many atoms. Also, in the low-densityenvironment of the ISM, collisions that would result in the change of a state, occur on timescales ofhundreds of years. From observations of that particular photons we learned [1], that H1 clouds have

†The term ’color excess’ is explained in the subsection 2.2.†Hydrogen comprises approximately 70% of the mass of matter and helium makes up for most of the rest; metals,

such as carbon and silicon, account for only a few percent.

7

temperatures of 30 to 80 K, number densities in the range of 1× 108 m−3 to 8× 108 m−3 and masseson the order of 1 − 100M�. More on the subject of observation of absorption lines of hydrogen issaid in the seminar of Vid Irsic (Lyman-α prehod v spektrih oddaljenih kvazarjev, 2010).

p/pmax

λ/λmax

(a) (b)

Figure 5: (a) Wavelength dependant polarization of light passing through the dust cloud. [5] (b) Fine andhyperfine hydrogen structure.[6]

Comparing observations of hydrogen optical depth τH and extinction in visual band AV alongthe same line of sight generally suggests that dust and gas are distributed together throughout theISM. However, sometimes column density of H1 is no longer proportional to column density of dust.That happens when dust becomes optically thick. These clouds can shield hydrogen from ultravioletsources and consequently molecular hydrogen can exist without udergoing dissociation by UV photonabsorption. Dust can also enhance the H2 formation rate. If the column density of hydrogen issufficiently large (NH on the order of 1025 m−2), it can also shield H2 from UV photodissociation.Molecular clouds are surrounded by shells of H1. The detection of H2 is difficult since it has noabsorption or emission lines in radio or visible part of EM spectrum at the cool temperatures typicalof ISM. In most instances it becomes necessary to use other molecules as tracers of H2 by makingthe assumption that their abundances are proportional to the H2. The most commonly investigatedtracer is carbon monoxide.

3.2 Milky Way, LMC nad SMC extinction curve

In this section, we will look at main properties of three galaxies (Milky Way, Large MagellanicCloud and Small Magellanic Cloud). Those three have significantly different properties and can beused as models for different types of galaxies (when discussing extinction or dust-to-gas ratio). Anondimensional form of the empirical extinction curve can be defined as [4]

ξ (λ) =AλAB

, (3.12)

where AB refers to the extinction in magnitudes in blue photometric band. The results of interstellarextinction measurements are commonly analyzed using the color excess Eλ−V ≡ Aλ − AV , dividedby EB−V , so

ξ (λ) =(Eλ−V /EB−V ) +RV

1 +RV

, (3.13)

where RV ≡ AV /EB−V is the ratio of total-to-selective extinction and the subscript V refers to thevisual photometric band. RV is independant of λ and its value can be determined by extrapolation

8

of ξ (λ) to infinite wavelength [4]. One can assume that ξ (λ) → 0 as λ → ∞, implying RV =−E∞−V /EB−V from equation (3.13). The extrapolation method gives more accurate values of RV

than other methods that attempt to measure AV directly†. The accuracy depends on the assumedform of ξ (λ) at large wavelengths∗.

0

1

2

3

4

5

6

0 2 4 6 8 10

ξ(λ)

λ-1 (µm

-1)

MWLMCSMC

Figure 6: Mean extinction curves for MW, LMC and SMC. [4]

Extinction is generally measured toward a pair of reddened and unreddened stars with nearlyidentical spectra. In the Milky Way, the O or B main-sequence stars are chosen for comparison,because they have distinctive spectra and thus can be identified with high precision, and they arealso highly luminous and thus can be seen to large distances. Direct measurements have been donefor Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC) as well, but because of thelarge distance from us (55 and 56 kpc), OB supergiants have to be used for extinction measurement.For any particular star, the observed values of Eλ−V and EB−V are proportional to the number ofdust grains along the line of sight (see (3.11)), which varies from star to star. Normalisation of Eλ−Vwith EB−V gives us the average extinction curve from measurements along different lines of sight.Empirical curves for MW, LMC and SMC are shown in Figure 6.

The three curves are similar at the infrared and visual bands, but differ substantially in theultraviolet wavelengths. The far ultraviolet extinction is smallest for the Milky Way and largestfor SMC. On the Milky Way curve is a remarkable feature: the broad bump centered near 2175 Awith an average width (FWHM) of about 490 A. The existence of the bump gives us some hint ofthe composition of the dust. For example, graphite interacts strongly with light near 2175 A. Theexistence of the bump, it’s strength and the abundance of carbon have led most researchers to suggestthat graphite may be a major component of interstellar dust. Another possible source of 2175 Afeature may be polycyclic aromatic hydrocarbons (PAHs), complex organic molecules with multiplebenzene structures. Not visible in Figure 6 are also contributions to extinction at wavelengths9.7 µm and 18 µm in the near-infrared. These features are believed to be the result of stretching ofthe Si-O molecular bond and the bending of Si−O − Si bonds in silicates.

†One of the possible methods uses the correlation, shown in Figure 1 left.∗For example, in [4], the assumed form is ξ (λ) = a0

1+(λ/λB)2.

9

EB

-V

NHI / 1021 cm-2

MW

1/λ (μm)

ξ(λ)

Figure 7: Left: correlations between color excess and column density for MW, LMC and SMC. Right:comparison between the model and three empirical extinction curves. Relative contributions from graphiteand silicate grains are shown (dashed lines), with the sum of the two shown as the solid lines.[4]

Next we can consider dimensionless dust-to-gas ratio [4]

k ≡ 1021 τBNH1

cm−2, (3.14)

where τB is the optical depth in the B band and NH1 the column density of neutral atomic hydrogen.Using (3.10) to express τB with EB−V and definition of RV , we can write dust-to-gas ratio as

k = 9.21× 1020 (1 +RV )EB−VNH1

cm−2. (3.15)

Here we discuss only neutral atomic hydrogen because there has been little information about molec-ular hydrogen on scales of individual clouds in LMC and SMC. In Figure 7 correlations betweenthe color excess EB−V and H1 column density for all three galaxies are shown. The straight linesrepresent the best fits of the ratio EB−V /NH1

†. The three mean ratios are significantly different.Equation (3.15) gives the value k = 0.78, 0.16 and 0.08 for MW, LMC and SMC, respectively. Em-pirically, the decreasing of ratios appears to be correlated to the decreasing trend of the heavy metalabundances.

Using the observations one can build best fitted model for the extinction, taking different typesof grains into the account. In the case of our three types of galaxies, only two major componentsof grains are necessary to include in the consideration: graphite and silicate grains. Grain sizedistribution can be approximated with a simple power-law of the form N(a) ∝ a−β, with β ≈ 3.5[4]. With that we can calculate the form of extinction curve as well as the contribution of specificgrains. The results are given in Figure 7.

Based on a good model, fitting function for the extinction curves can be built. As mentionedbefore, the three mentioned models can be used as general models for different galaxies, as theyare quite distinct. However, those models are not the only ones possible [7]. Other models can beconstructed with different grain size distributions and compositions, models describing supernovaevent. . .†The measurement of column density will be discussed in later chapters.

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3.3 Dust destruction

One of the important issues of extinction in GRB host environment is destruction of the dust. Inprincipal, two possible mechanisms are important for dust destruction: grain heating and graincharging [8]. The former happens when the grains absorb energy fasterthan they manage to radiateaway. The excess energy than goes into breaking the bonds that hold atoms to the surface of thegrains. The charging can lead to electrostatic stresses greater than the tensile strength of the grains.Typical distances, at which the destruction has an effect, are 1018 − 1020 cm. The destructionresults in lesser obscuration of the optical/UV afterglow. Also, additional absorption lines of heavierelements, returned to the gas phase, can be observed in the afterglow spectrum. Since the radiation atearlier time after the burst destroys the surrounding dust region, the spectral energy distribution andconsequently the physical properties deduced from the afterglow can differ substantially at differentepochs after the burst. For example, the reddening AV at the start of the afterglow radiation mightdisappear at later epochs.

Figure 8: The gradual recycling of carbon from dust into gas during the passage of the radiation flux.[8]

During the destruction, the fraction of the heavier elements in the gas will increase. In Figure 8such increase of carbon in shown in simulation. That results in the change of the extinction curve.Since the curves (especially that of Milky Way) have a strong dependance on carbon, they will nowexperience the change, shown in Figure 9. The curve becomes narrower at ultraviolet frequenciesand the bump dissapears.

4 Extinction in GRB host galaxies

The study of GRB afterglows is vital in the process of obtaining the values of intrinsic light extinctionand NH of GRB host galaxies. Since the empiric light curves studied in section 3.2 only describe thewavelengths from UV to IR band, the higher energy X-ray bands have to be studied independently.Measuring the amount of the X-ray extinction yields information on the total NH from the observerto the source. Although the amount of the X-ray extinction is usually expressed in terms of NH , itis generally caused, especially above ∼ 0.25 keV, by the most abundant heavier elements like O, Ne,Fe, Mg and Si [10]. One usually obtain NH by fitting the measured X-ray spectrum from the source

11

Figure 9: The evolution of the extinction curve, as the radiation passes through it, altering the graindistribution. [8]

with models, constructed from individual components. For example, a simple model would includea source (power law or broken power law; the exact form in both cases is treated later), Galacticabsorption and host galaxy absorption. The model for (photoelectric) absorption can be taken as[11]

e−NHσ(E), (4.1)

where σE is the cross-section. That is actually the exponential decrease of light because of absorp-tion (see chapter 3). Based on the speculation and knowledge already obtained, one can use theappropriate form of the cross section (i.e. photo-electric Wisconsin cross-section [12]) with one ofthe already calculated relative abundances set. This model was used as the best fit in the examplein Figure 10 (left).

The light extinction AV is obtained by fitting [7]

Fν = F0(ν)10−0.4A(ν), (4.2)

where F0(ν) is the afterglow spectrum in the absence of extinction: F0(ν) ∝ ν−βo,i (βx = βo,i) in thepower-law case, F0(ν) ∝ ν−βo,i (for ν < ν0) and F0(ν) ∝ ν−βx (βx = βo,i + 0.5 †, for ν > ν0) in thebroken power-law case (ν0 is the break frequency). A(ν) is the extinction profile and depends on thedust model assumed (section 3.2).

The example of the whole fit is shown in Figure 10 (right). The best fitting extinction profile inthis case is that of SMC profile with the value of extinction AV = 0.48 ± 0.19 mag. The hydrogencolumn density derived from X-ray spectrum is NH = (5.8± 0.4)× 1021cm−2.

In Figure 11 additional examples of various spectral fits are shown. Although only MW, LMCand SMC extinction profiles are shown, others may also be considered (models with different grainsize distribution or composition, supernova models...) [7].

In [14], a thourough investigation of a sample of 28 GRB spectral energy distributions have beenmade. It has been found that 79 per cent of the sample has a host galaxy absorption system detectedwith an equivalent neutral hydrogen density ranging from NH = 8.2 × 1020 to 1.4 × 1022cm−2. Adust-extinction system local to the GRB is detected in 71 per cent of the sample with a range in

†The theoretical constraint.

12

10.5 2 50.01

0.1

110

norm

aliz

ed c

ount

s s−

1 ke

V−

1

Energy (keV)

GRB061007 XRT

Figure 10: Left: absorbed power law fit of GRB061007 X-ray part of the spectrum. Right: optical to γ-rayspectral energy distribution, fitted with an absorbed power law. The unabsorbed power law is also shown(dotted line). [13]

F(m

Jy)

Frequency (Hz)

Figure 11: Best-fitting absorption and extinction corrected spectra (solid lines), as well as the host galaxyabsorbed spectral model (dashed lines) and data (open data points) for three GRBs.[14]

extinction of 0.03 < AV < 0.75. The distribution of host galaxy column densities is shown in figure12 (left; solid line). Also, the expected NH distribution is shown (dashed line) in the case of Galactic-like molecular clouds. Both distributions would coincide in the case of solar metallicity GRB hostgalaxies. From measurements so far, the GRB host galaxies typically have subsolar metallicities,ranging from 1/100th of solar to solar [14]. However, due to the small number of accurate metallicitymeasurements, the metallicity distribution is poorly known. Figure 12 also shows measured opticalextinction distribution with a mean value of AV = 0.3 and standard deviation of 0.2. It was alsofound that the SMC extinction curve provided the best fit in majority of cases. Since SMC extinctionprofile has no 2175 A bump, the result suggests that small carbonaceous dust grains are not typicalin the environments of GRBs, at least not once the GRB has occured.

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Figure 12: The measured distribution of column density (left) and optical extinction (right).[14]

4.1 The highest redshift GRB

GRB 090423 was detected with Swift on 23 April 2009 [15]. It turned out that the GRB originatedat redshift z = 8.3 - the most distant object ever observed. Beyond breaking a distance record, thehigh redshift has several crucial implications. First, it demonstrates that massive stars did indeedexist at that time (∼ 625 million years after the Big Bang). Second, it can prove invaluable to thescience of cosmology, since it is close to the redshift range during the bulk of cosmic reionization isthought to have taken place. Also, future deep observations of the host galaxy may provide strongerconstraints on the nature of the underlying stellar population, and hence the possibility that GRB090423 originated in a galaxy undergoing its first burst of star formation.

The spectral energy distribution of the burst is shown in Figure 13 at two different times. Thetwo spectral fits show the time evolution. As it has been suggested in the chapter 2.1, the breakingpoints are moving toward lower frequencies and that is clearly visible at that example. It mustbe emphasized though, that such spectral evolution is present in many observed spectra and is notrelated to the distance of this particular burst.

The early time SED suggests best fit values βx = 0.80 with βIR = 0.30 (with a fixed ∆β = 0.5).No host extinction fit is required. The late time fit shows that the X-ray to optical fit has increased toβX−IR = 0.90. The best-fitting model of X-ray SED suggests column density NH = 1.16×1023 cm−2.

14

Figure 13: SED of the GRB 090423, measured at different times. The left hand panel shows only theoptical/IR part of the spectrum, while the right panel includes the XRT observations.[15]

5 Conclusion

Since the launch of the Swift satellite and consequently fast response times of GRB afterglow detec-tion, the research of GRB host galaxy properties have become much easier. The dust compositionand hydrogen column density obtained with the help of the discussed analysis allow us to look atthe distant galaxies and compare the galaxy properties at different times since the Big Bang. Therecent results suggest the lack of carbon in the gama-ray burst galaxies. The argument against thatconclusion could provide the dust destruction, since the simulations show the decrease of carbon inthe grain form. While that could be true, the small fraction of LMC curves that best describe theextinction present us with counter argument. Also, the small average value of extinction obtainedin the analysis present a problem yet to be explained. It can be expected, that in near future manymore GRB events will be thouroughly analysed, thus providing better understanding of the galacticevolution as well as the GRB events itself.

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References

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