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Gamma-Ray Burst afterglows
Galama, T.J.
Publication date 1999
Link to publication
Citation for published version (APA): Galama, T. J. (1999). Gamma-Ray Burst afterglows.
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Download date:26 Sep 2021 Introduction Introduction
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Introduction n
1.11 What are 7-ray bursts? Gamma-rayy bursts (GRBs) are the strongest phenomenon seen at 7-ray wavelengths; bright GRBss are as bright in 7 rays as the brightest star in the sky, Sirius, is in visible light. GRBs weree discovered with the Vela satellites, whose main purpose was to verify compliance with the 19633 Limited Nuclear Test Ban Treaty. Since their discovery (Klebesadel, Strong and Olson 1973)) these events, which emit the bulk of their energy in the 0.1 — 1.0 MeV range, and whose durationss span milliseconds to tens of minutes, posed one of the great unsolved problems in astrophysics.. Until recently, no counterparts (quiescent as well as transient) could be found andd observations did not provide a direct measurement of their distance. The breakthrough camee in early 1997, when the Wide Field Cameras aboard the Italian-Dutch BeppoSAX satellite allowedd rapid and accurate localization of GRBs. Follow-up on these positions resulted in the discoveryy of X-ray, optical and radio afterglows. These observations revealed that GRBs come fromm 'cosmological' distances, and that they are by far the most luminous photon sources in thee Universe, with peak luminosities in 7 rays up to 1052 erg/s, and total energy budgets up too several times 1053-54 erg (for assumed isotropic emission). The optical signal from GRB is regularlyy seen to be 10 magnitudes brighter (absolute) than the brightest supernovae, and once evenn 18 magnitudes brighter.
1.1.11 Vela satellites Thee Limited Nuclear Test Ban Treaty prohibits nuclear weapons tests "or any other nuclear explosion"" in the atmosphere, in outer space, and underr water. The Vela satellites were designed suchh that they could verify compliance with the treaty by detecting the 7 rays from nuclear testss outside the Earth's atmosphere. Vela 5A, 5B, 6A and 6B each carried six 10 cm3 Csl scintillationn counters; they could detect photons in the 0.2-1.0 MeV (Vela 5) and 0.3-1.5 MeV (Velaa 6) energy range (see Fig. 1.1). AA search for 7-ray bursts was started by R. Klebesabel because of the prediction that 7-ray emissionn would be observable during the initial stages of supernova explosions (Colgate 1968). Noo indications for such a phenomenon were found. It was in 1969 that, embedded in Vela spacecraftt data from 1967, a 7-ray burst was found. At that time the Sun could not be exluded ass a source. With the launch of a new generation of Vela satellites (the Vela 6) sufficient timing accuracyy made it possible to exclude the Sun as the source of these events. Once about 16 eventss were found (and believed) a paper was published announcing the discovery of cosmic
7 7 ChapterChapter I
7-rayy bursts (Klebesadel, Strong and Olson 1973). The original July 1967 event is not in that paperr because it could have come from the Sun (although its characteristics are like a GRB and itt is now considered the oldest known GRB)'. A time history of the 1967 event can be found in Strongg and Klebesabel (1976).
Figuree 1.1: Left figure: Vela-SA and 5B Satellites in the Clean Room. Right figure: animation ofof Vela-5B in low Earth orbit (courtesy of Los Alamos National Laboratory).
1.1.22 Light curves and spectra of GRBs Thee Burst And Transient Source Experiment (BATSE; Fishman et al. 1989) on board the Comp- tonn Gamma-Ray Observatory (CGRO; see Fig. 1.3) observes about one GRB a day. The CGRO providess a wealth of information on GRB light curves and spectra; here I summarize some of thee most important results. A more extensive discussion on temporal properties and spectra of GRBss can be found in the review by Fishman and Meegan (1995).
GRBs have durations ranging from milliseconds to ~ 103 s. Their time histories display aa great diversity of structure, with single or multiple peaks, with smooth profiles, profiles withh sub-peaks, profiles with well separated peaks and profiles with overlapping peaks andd spikes (see Fig. 1.2).
GRB light curves show rapid variability, on time scales sometimes less than a millisecond (Bhatt 1992; Schaefer and Walker 1999). As light travels a distance L = ct = 3-107cmin onee millisecond, about 10 neutron star radii, it is generally believed that GRBs originate fromm compact objects, such as neutron stars (NS) or black holes (BH).
Several attempts have been made to categorize GRB time histories. No other observa- tionall parameters appear to be correlated with temporal morphologies (e.g., Fishman and Meegann 1995), i.e., there is little morphological evidence for distinct burst classes within
11 It is often believed that the publication was delayed until 1973 because the data were classified. In fact it was thee better timins of GRBs which was needed to rule out the Sun.
8 8 Introduction Introduction
thee GRBs. There is one exception: the distribution of burst durations is bimodal and separatess GRBs into two classes, the short events (<2s) and the longer ones (>2s) (Kou- veliotouu et al. 1993; see Fig. 1.4). The duration appears to be anticorrelated with spectral hardness:: short bursts are predominantly harder than long ones (Kouveliotou et al. 1993).
Norriss et al. (1995) find evidence for time dilation by comparison of samples of bright andd dim BATSE GRBs; the centroids and widths of the duration distribution for the dim samplee are scaled by a factor of two relative to the bright sample. They interpret this as a resultt of cosmological redshift: the dim bursts are, on average, located at larger distance (i.e.,, redshifts) than the bright bursts. The sources of dimmer bursts would lie at redshifts off order 2.
High-energyy emission is a unique feature of GRBs. Spectral measurements extend from a feww keV to ~ GeV (see e.g., Fig. 1.5). The continuum spectra of GRBs are very broad and hard;; most of the power is emitted above 50 keV. GRB spectra are well described by an empiricall function, the so called Band function (Band et al. 1993). This function consists off a low- and high-energy power law, smoothly joined by an exponential turnover. It has threee parameters, the peak energy, and the low- and high-energy photon index (for details seee Bandetal. 1993).
Delayedd very high-energy 7 rays have been observed from GRBs with EGRET onboard thee CGRO (Hurley et al. 1994; Sommer et al. 1994). In the case of GRB 940217, 200 MeV-200 GeV photons were observed with EGRET up to 90 minutes after the event onset (Hurleyy et al. 1994), i.e., much longer than the duration of the lower photon energies of thee GRB itself (~ 200 seconds; 30-2000 keV).
Fordd et al. (1995) investigated the evolution of the peak energy for long and bright GRBs andd found that the peak energy decreases with time. Liang and Kargatis (1996) found thatt the peak energy decreases exponentially with the photon fluence. This 'hard-to-soft' spectrall evolution (see also Norris et al. 1986) is also found within individual pulses of GRBss (Ford et al. 1995).
Timee histories of GRBs are different for different energy bands. At higher energies the overalll burst duration as well as the rise and fall time scales of pulses are shorter than thosee at lower energies (Link, Epstein and Priedhorsky 1994). Typically, the low-energy emissionn in GRBs persists longer than the high-energy emission.
1.1.33 Distribution on the sky and in space Priorr to the launch of the CGRO, in 1991, it was found, from various experiments, that the burstt population is approximately uniform in space. Detector sensitivities were still too small to detectt an 'edge' to the distribution of burst sources (Higdon, Meegan and Cline 1984; Hurley, Clinee and Epstein 1984). Thee space distribution of GRB sources is related to their apparent flux distribution, as can bee seen from the following simplified argument. We assume that bursts are standard candles, i.e.,, they emit the same amount of luminosity at the same wavelengths. A burst at a distance 2 dodo is observed with a peak flux P0, P0 ^c d^ . The cumulative rate of bursts, with peak fluxes
9 9 ChapterChapter 1
Burstt number : : 105 5 Burstt number: 249 bU U -- 50 0 fl l 40 0r r -. -. 30 0r r -_ -_ 20 0 -- 10 0 0 0 ... i... i .. i . . : : -55 0 5 10 15 -100 0 10 20 30 40 50 60 Timee (seconds since trigger) Timee (seconds since trigger)
Burstt number: 257 Burs B t t number: : 408 8 14 4 ...,...,...,...,.. . -- 12 2 --
10 0~ ~ -- ^JU diU di 8 8 yn^ww yn^ww fi fi , , -100 0 10 20 30 40 50 -200 0 20 40 60 80 100 120 Timee (seconds since trigger) Timee (seconds since trigger)
Burstt numt er:: 677 Burstt number: 678 5U U 40 0 --
30 0 --
20 0 -_ -_
10 0 ^i^l/^AfV^!^ ^i^l/^AfV^!^
n n-- i -1.00 -0.5 0.0 0.5 1.0 00 20 40 60 Timee (seconds since trigger) Timee (seconds since trigger)
Figuree 1.2: A sample of GRB light curves (25-2000 keV) observed by BATSE. The light curves displaydisplay a large variety in time profiles, duration, and intensities. Trigger 105 is an example ofof a burst with smooth, well defined peaks. Trigger 257 is a typical example of a Fast Rise, ExponentialExponential Decay (FRED) light curve. Trigger 408 is an example of a complex, chaotic and spikyspiky burst. Trigger 677 is a GRB with very short duration, and belongs to the sub-class of shortshort GRBs (< 2 sec) (courtesy of the BATSE team).
PP exceeding P0, R(> Po), is proportional to the volume \0 of space in which bursts with ! 2 observedd peak fuxes P > P0 can be observed. Hence it is proportional to rf0 and, by P0 ~x d^ , i/2 proportionall to P0T' , yielding a log7?(> P0)-\ogP0 curve with a slope of-3/2; this is true evenn if the luminosities have a broad distribution. In the 1980's it was observed that bright burstss follow this slope of-3/2. However, balloon experiments, with very large area detectors, foundd that the weaker bursts deviate from the -3/2 slope (Meegan, Fishman and Wilson 1985). Relativelyy fewer faint sources than expected from a homogeneous distribution of bursts were seen.. It was not clear at that time whether this deviation was the result of selection biases or a reallyy existing effect.
10 0 Introduction Introduction
Figuree 1.3: The Compton Gamma Ray Observatory (CGRO) above western Africa just prior to itsits release from the Space Shuttle into orbit. Visible are the solar panels and the four CGRO experiments:experiments: the large round domes of the EGRET (bottom) and COMPTEL (center) experi- mentsments and four of the eight BATSE detectors located at the corners of the satellite are visible. TheThe OSSE experiment housing is visible just above the COMPTEL dome (courtesy of NASA).
BATSEE 4B Cotolog
Figuree 1.4: Distribution of T90 far BATSE burstsbursts from the 4B catalogue (Paciesas et
al.al. 1999). T90 is defined as the time during whichwhich the cumulative counts increase from 5%5% to 95% (Kouveliotou et al. 1993).
(seconds) )
Thee BATSE surprise (1991)
Astronomicall objects in the vicinity of the Sun show a uniform distribution in space and are isotropicc on the sky. The general notion in the GRB community, based on these observational
11 1 ChapterChapter 1
Miiitii ; iii'in; i :imiii i \\\ PIJ i ii i rrmi ' i inn—!—I i linn—r > > w w -GRB910503 3 GRB910814 4 .5 5 o o o o o o M M o o /^k /^k A A 2 2 ? ? i i > > > > l; ; .1 1 4-- -* i i thi i I I -A -A tt • 't—• "" t fa fa Of, Of, 4 j4 j t t 1 1T T .022 [—/ c c I I JZ JZ
LLUllLLUll l,..L,i. lit ttl\ III ill ai I ll III,ill l 100 10; 10' lüs itrr to' Energyy (keV) Energyy (keV)
Figuree 1.5: 5pecfra of two bright GRBs (GRB 910503, 20 keV to 300 MeV; GRB 910814, 103 keVkeV to 500 MeV) from BATSE, COMPTEL and OSSE data (from Schaefer et al. 1998). facts,, was that GRBs were produced by neutron stars and that we observed the nearby ones inn the solar neighborhood. This notion was reinforced by the alleged detection of cyclotron liness (reflecting a magnetic field ~ 1012 Gauss) in the spectra of several GRBs with the Venera andd Ginga satellites (Mazets et al. 1981; Murakami et al. 1988; Fenimore 1988), and by emissionn features around 400 keV, which were interpreted as gravitationally redshifted 511 keV annihilationn lines (Mazets et al. 1981). However, in spite of very extensive searches, BATSE hass not found any such spectral features (Palmer et al. 1994; Band et al. 1996). Itt was believed that the Milky Way of bursters would be revealed to a more sensitive BATSE. Thee big BATSE surprise (Meegan et al. 1992) was that even faint GRBs are distributed isotropi- callyy on the sky (see Fig. 1.6). In addition, the fainter bursts showed a turnover in the cumulative brightnesss distribution (see Fig. 1.7), i.e., there is a distinct dearth of very weak GRBs. BATSE alsoo sees no clustering of bursts on small or large angular scales, i.e., GRBs are not associated withh concentrations of mass on any distance scale (e.g the Galactic disc, nearby clusters of stars,, the Large Magellanic Cloud, nearby galaxies like M31 or clusters of galaxies like Virgo). Thee simplest explanation for these observations is that we are at the center of a spherically symmetricc distribution of GRBs and that we are observing the 'edge' of this distribution.
1.1.44 The 'great' debate
Thee BATSE observations (see the previous Sect. 1.1.3) excluded that GRBs come from a galac• ticc disk population (Meegan et al. 1992); a Galactic disc population of burst sources would ei• therr be homogeneous, at distance scales less than several hundred pc, or anisotropic, at distances greaterr than several hundred pc. Naturally these results led to the conclusion that GRBs orig-
12 2 Introduction Introduction
23000 BATSE Gamma-Ray Bursts
Figuree 1.6: Angular distribution of BATSE bursts in Galactic coordinates. inatee from 'cosmological' (Gpc) distances (Paczyriski 1986, 1992, 1995; Mao and Paczynski 1992a).. There was, however, a countervailing view that GRBs originate in a very large halo aroundd our galaxy (Brainerd 1992; Lamb 1995; Mao and Paczynski 1992b). As a result, the discussionn on the nature of GRB sources focused on their distance scale.
Galacticc halo
Alll known Galactic objects are strongly concentrated to the Galactic center. In order to have noo measurable anisotropy of a Galactic population of bursts, the objects should populate an extendedd halo with a very large core radius, R > 100 kpc, so that the offset of the Sun from thee center of the Galaxy does not show up. The lack of any observable concentration of bursts towardss the nearby galaxy M31 gives some further constraints on the size of a Galactic burster halo.. The halo radius was estimated to lie between 100 to 300 kpc (Briggs et al. 1996; Hartmann etal.. 1995; Hakkilaet al. 1994). This would imply that bursts should populate a yet unobserved veryy extended hypothetical halo. It was suggested that such an extended halo might perhaps bee created by injection of high-velocity neutron stars from the Galactic plane, escaping the gravitationall field of the galaxy (e.g., Hartmann 1993). Assuming a distance d = 100 kpc for 42 aa Galactic burst we find a luminosity LGRB ~ 10 ergs~' (for a typical GRB flux, fGmi = 10"6ergcm~2s-1,, and assuming isotropic emission).
Cosmological l
Ass far as we know, on much larger distance scales, several Gpc, all objects are distributed roughlyy isotropically and uniformly. A cosmological distribution of sources would reveal a naturall deficiency of weaker bursts by relativistic effects that affect the weaker (on average att greater distances) bursts (Paczynski 1986, 1992; Mao and Paczynski 1992a). Because of cosmologicall time dilation the count rate is lower by a factor 1 + z. Also, space is not Euclidian,
13 3 ChapterChapter 1
-2-100 1 2 -2 log100 P (ph cm s~' 50-300 keV)
Figuree 1.7: The cumulative peak flux distribution of "triggered" BATSE bursts combined with thethe "non-triggered" BATSE bursts found in 6 years of archival BATSE data (solid line; units areare bursts yr~l sr~]; from Kommers et al. 1999). Shown are the best-fit cosmological mod- elsels with power-law luminosity distributions for: a co-moving burst rate that is independent of redshiftredshift (dot-dashed line), and for co-moving burst rates that follow the star-formation rate as determineddetermined by Madau et al. (1998) and Hughes et al. (1998). The slope — 3/2 for a uniform distributiondistribution of sources is also indicated. whichh becomes noticeable at larger scales, and the observed spectrum is redshifted by a factor 11 + z, i.e., a different part of the spectrum is observed (e.g., Piran 1996). For a typical GRB flux,flux, /GRB = 10-6ergcm~2s~', and a cosmological distance of d = 3 Gpc we find a luminosity LoRB-KFergs-1. .
1.1.55 Counterparts
Thee lack of knowledge about the nature of GRBs led to a lot of speculation and an enormous numberr of proposed models. Some examples are the merger of a double neutron star (NS-NS) orr of a neutron star and black hole (NS-BH) binary, failed supernovae, asteroids or comets fallingg into black holes or neutron stars, processes in the core of active Galactic nuclei etc. (see Nemirofff 1994, for a list of more than 100 models of GRBs published before 1992). Ass mentioned above, the 7-ray properties of GRBs did not provide unambiguous distance clues,, and it was generally agreed that settling the debate on the GRB origin required the iden- tificationn of GRB counterparts at other wavelengths (e.g., Schaefer 1994). Identification of aa known type of object associated with a GRB would immediately reveal the distance scales andd thereby greatly constrain theoretical models of GRB production. However, the generally
14 4 Introduction Introduction largee positional errors of bursts and the fact that accurate positions were usually obtained only longg after the event, made searches for counterparts difficult. Historically, three strategies were employedd for counterpart searches. Deep searches for quiescent counterparts to accurately lo- calizedd events, i.e., low-energy emission long after the burst, have been made and reveal that quiescentt counterparts are very faint, i.e., not detected, at all energies (Schaefer 1994). For example,, in radio at 2, 6, and 20 cm, upper limits of ~ 100-800 (iSy on quiescent counterparts forr ten small GRBs error regions were given by Schaefer et al. (1989). Another strategy has beenn to search for flaring counterparts in simultaneous wide-field monitoring experiments in thee hope to have a GRB included in the field of view of the instrument at the moment of the event.. A third counterpart search technique is to use the GRB detection as a trigger to point the telescopee in the appropriate direction and search for a flaring and/or fading counterpart on time scaless as long as, and much longer than the burst itself; it is this third technique that proved to bee successful (see Sect. 1.3.1). Very recently also optical emission simultaneous with a GRB wass discovered (Akerlof et al. 1999).
1.22 Relativistic blast-wave models
Fireball-plus-relativisticc blast-wave models, first proposed by Goodman (1986) and Paczynski (1986),, predict a late-time low-energy radiation (see, e.g, Mészaros and Rees 1997; Paczynski andd Rhoads 1993), and, together with some more phenomological considerations, provided the incentivee for X-ray, optical and radio counterpart searches within hours or days after the event. Thee basic model is a point explosion with an energy of order 1052 ergs, which leads to a 'fireball',, an optically thick radiation-electron-positron plasma with initial energy much larger thann its rest mass (e.g., Piran 1996; see Piran 1999 for an extensive review) that expands ultra- relativistically.. The GRB may be due to a series of 'internal shocks' that develop in the rel- ativisticc ejecta before they collide with the ambient medium. When the fireball runs into the surroundingg medium a 'forward shock' ploughs into the medium and heats it, and a 'reverse shock'' does the same to the ejecta. As the forward shock is decelerated by increasing amounts off swept-up material it produces a slowly fading 'afterglow' of X rays, then ultraviolet, optical, infrared,, millimetre, and radio radiation (see Fig. 1.8). In the following sections I will discuss thiss model in more detail.
1.2.11 The compactness problem, relativistic motion and baryon loading Thee observed GRB spectra contain a large fraction of high-energy 7-ray photons (see Fig. 1.5). AA high-energy 7-ray photon (of energy E\) may interact with a lower-energy photon (of energy EE22),), to produce an electron-positron pair, when the two photons fulfil the condition
2 2 EEXXEE22{1{1 - cos012) > 2{mec ) , (1.1) wheree 012 is the angle between the directions of the two 7 rays, me is the rest mass of the electronn and c is the speed of light (e.g., Carrigan and Katz 1992). Thee observed rapid variability in GRB light curves implies that the sources are compact, RR < cST ~ 300 km, where ST is a variability time scale. The large distances for cosmological burstss imply that bursts release a large amount of energy (~ 1052 ergs). Under these circum- 11 stancess the optical depth to pair production, 77 -> e , is very large, r77 ~ 10 (e.g., Woods
15 5 ChapterChapter 1
£»» m jsr»ikin>i! ! "-:._ "-:._ tte&m tte&m \ \ ass Y'-rm% Optiotll \ urur UV, \ ««HJ5JJ tit's é&zags é&zags uss tiw l&fti
hiU'rmtihiU'rmti SI«H'fo Forwardd shert (•• Utrnal Sis>i<.k
Figure e .8:: Schematic overview of the fireball-plus- relativistic blast-wave model (courtesy Dr. Piran) Piran) andd Loeb 1995); many pairs will form and these will result in a huge optical depth for all pho• tons,, i.e., we expect a blackbody form of the radiation. However, the observed spectra are not thermal! ! Historically,, it was realized that relativistic motion can explain how the very high luminosity inn 7 rays can occur without 'self destruction' of the photons via electron-positron pair produc• tionn (- );); the photons with an observed energy huobs have been blueshifted and their energyy at the source was lower by one factor T, ~ fw0iK/T, where T is the Lorentz factor of thee outflow at the time of the 7-ray emission. Also, the radius where the radiation is emitted cann be larger by a factor T2, and relativistic motion beams the photons on the average within ann angle equal to T_1 (both effects reduce the optical depth). Since very energetic photons are observedd this indicates large bulk Lorentz factors, F > 102 (e.g., Fenimore, Epstein and Ho 1993;; Woods and Loeb 1995; Piran 1999). Only if the baryonic load is small will the outflow remainn relativistic. However,, astrophysical fireballs are contaminated with ordinary matter (baryons) and a con• versionn of radiation energy into kinetic energy will take place. Under most circumstances, even aa small baryonic load will cause the transfer of almost all of the fireball energy into kinetic energyy of the baryons (Shemi and Piran 1990; Paczynski 1990). A mechanism was needed too convert this energy back to radiation. Rees and Mészaros have therefore suggested that the observedd "-ray emission is due to a secondary phase of photon emission (Rees and Mészaros 1992,, Mészaros and Rees 1993), caused by the interaction of the relativistic expanding fireball andd the interstellar medium. The bulk kinetic energy of the baryons can be re-converted into heat,, and then radiated as 7 rays, in an external shock, when this highly relativistic blast wave iss slowed down by the ambient gas. Since this occurs at much larger radii, the fireball can be opticallyy thin and so the spectrum can be non-thermal (Rees and Mészaros 1992, Mészaros and Reess 1993). Similarly the GRB could be produced when the relativistic blast wave is slowed downn by self interaction in internal shocks (Rees and Mészaros 1994; Paczynski and Xu 1994). Sarii and Piran (1997) argue that the internal shock scenario for GRB production is more likely
16 6 Introduction Introduction too work than the external scenario, but this is not universally accepted (e.g., Dermer and Mitman 1999). .
1.2.22 Blast-wave dynamics Followingg the production of the GRB an external shock is expected as the fireball interacts with thee surrounding inter-stellar medium gas (ISM). This second interaction (or first, depending on whatt produced the GRB) then accounts for the observed afterglow at X-ray, optical and radio wavelengths.. We use the term interstellar medium but note that the gas need not necessary bee interstellar. A cold shell is formed that expands with approximately uniform Lorentz factor rr = 77. The interaction between the outward moving shell and the ISM takes place in the form of twoo shocks. A forward shock that propagates into the ISM and a reverse shock that propagates backk into the relativistic shell. We have therefore four distinct regions: the cold (unshocked) shell,, the shocked shell, the shocked ISM and the ISM (cold and unshocked). Thee relevant jump conditions across the shock follow from the continuity of energy, momen- tumm and particle flux densities (Blandford and McKee 1976; Sari, Narayan and Piran 1996). For aa strong shock we have
n'n'22 = 4Tn! (1.2) 7ii = v^T, (1.3) where:: the subscript 1 denotes the cold unshocked ISM, the subscript 2 denotes the shocked ISM,, primes denote comoving-frame quantities, n\t2 are the particle densities, 712 = y/UÖ^Ïï) aree the Lorentz factors, and 72 = T. The contact discontinuity (the boundary between the shockedd shell and the shocked ISM) moves with a Lorentz factor F and the forward shock with y/2Ty/2T (Eq. 1.3). The width, Ar, of the shocked ISM shell is therefore Arr = Af3ct = {\/AT2)ct = (l/4r2)r, (1.4) wheree r — ct is the shell radius. In the comoving frame an observer will see the ISM approach- 2 ingg with Lorentz factor T. The dominant kinetic energy is in the nucleons, Tmpc . To good approximationn all this energy is thermalized in a strong shock and hence the comoving energy density y 2 2 ee22 = n2E[herm = 4rc]r mpc , (1.5) s tne wheree E'therm i thermal energy per particle, and where we have used Eq. 1.2. The total energyy contained in the shell 2 E'E'tottot = 47rr Ar'e2. (1.6) Wee observe this energy blueshifted, hence
3 2 EEtottot = 47rr rVmpc - 3A/sweptrV, (1.7) wheree Mswept is the swept-up ISM mass. We can obtain an estimate of the deceleration radius, rdec,, the point where the energy in the hot, swept-up interstellar material equals that in the orig- inall explosion, by assuming that at this early stage deceleration does not yet strongly influence thee evolution of the blast wave and so T = rj = constant. Here r\ = E/MQC2 is the ratio of initiall energy in the explosion to the rest mass energy of the baryons (mass MQ) entrained in it. Wee derive 2 2 1/3 16 1 3 2 /3 rdecc = (£747T77 nmpc ) = 1.81 x 10 (£;g2/n) / 7?3-0 0 cm, (1.8)
17 7 ChapterChapter 1
Figuree 1.9: Synchrotron spectrum of a relativistic shock with a power-law electron distribution forfor the case of slow cooling, which is expected at late times. The spectrum consists of four segmentssegments (A-D). The frequencies //a, vm and, vc (see Sect. 1.2.3) decrease with time as indicated (from(from Sari et al. 1998).
2 wheree £52 = £/TCP erg and r/:SOo = ///300. The point of deceleration (Eq. 1.8) corresponds too Mswept = (I/377) A/Q, i.e., we have the counter-intuitive result that for a more energetic blast wavee (greater 77) less mass needs to be swept up to decelerate it. We can relate Eq. 1.8 to the timee in the observer's frame, t,obs, by
Sr Sr 5iobSS = ^2-, (1-9) andd using T = ?/ = constant, we find for the deceleration time scale
2 1/3 8 /3 *decc = rdec/2»? c = 3.35(JB52/n) r?30 0 s. (1.10)
Thee time scale of interaction of the blast wave with the ISM is of the order of seconds and it occurss at a rather large radius of 10lb cm, i.e., if GRBs are produced by the interaction of the blastt wave with the ISM, the radiation can be non-thermal. For internal shocks a similar size self-interactionn radius can be derived (e.g., Piran 1999). Here we will discuss this model for its predictionn of GRB afterglow.
Forr the case of an adiabatically expanding blast wave, we have Etot = constant, and so the blastt wave evolves as rr oc r":1/2. (1.11)
Usingg Eq. 1.9 we find
1/4 rr = rdec(tobs/*dec) (1-12) 3 TT = T](tohs/tdecr '*. , (1.13)
18 8 Introduction Introduction
1.2.33 Low-energy emission at late times: afterglow Fireballl shock models, in which relativistic electrons radiate via the synchrotron mechanism, predictedd that detectable optical and radio emission (Mészaros and Rees 1997; Paczyiiski and Rhoadss 1993) can be seen over a period of days to weeks after the GRB event. This late-time softerr radiation was dubbed the 'afterglow'. Thee observational characteristics of GRB afterglows can be derived from the classical syn- chrotronn spectrum (e.g., Rybicki and Lightman 1979), and taking into account synchrotron coolingg at high frequencies and synchrotron self absorption at low frequencies (e.g., Sari, Piran andd Narayan 1998). Thee electrons are assumed to be accelerated, in a strong shock, to a power-law distribution p off electron Lorentz factors, 7V(7e) oc 7e" , with some minimum Lorentz factor 7m. In the followingg we assume slow cooling (the bulk of the electrons do not radiate a significant fraction off their own energy and the evolution is adiabatic); these conditions appear applicable to some observedd GRB afterglows at late times. Then, the synchrotron spectrum of such a distribution llz off electrons is a power law with Fv oc v up to a maximum, Fm, at the peak frequency vvmm (corresponding to the minimum Lorentz factor 7m). Above vm it is a power law, Fv oc -1 2 2 2 iz-fp )/ ,, up to the cooling frequency, vc. Electrons with energies 7emec > 7cmec , where 7CC is the electron Lorentz factor associated with the cooling frequency uc, radiate a significant fractionn of their energy and thereby deplete the high-energy part of the electron distribution and,, correspondingly, the synchrotron spectrum above uc. Cooling causes a spectral transition; p 2 abovee uc we have Fv oc v~ / . Synchrotron self absorption causes a steep cutoff of the spectrum 2 att low frequencies, v < vA {Fu oc v if i/a < vm), where i/a is the synchrotron self absorption frequency.. Thus, the spectrum consists of four distinct power-law regimes, seperated by three breakk frequencies: (i) the self absorption frequency, ua, (ii) the peak frequency, vm, and (iii) the coolingg frequency, vc (see Fig. 1.9). Forr the synchrotron peak frequency in the comoving frame we have (e.g., Rybicki and Light- mann 1979) "m^^Tm,, (1-14) wheree B'2 is the magnetic field strength in the shocked ISM shell. We now assume that a constant fraction,, ce> of the shock energy goes into the electrons. Then (e.g., Wijers and Galama 1999, seee Chapter 10)
^^ = TTY — ^«er = «1r, (i.i5) 11 + A me p — 1 wheree X is the usual hydrogen mass fraction, and rae is the mass of the electron. Also we 2 assumee that the magnetic energy density B'2 /8-K is a constant fraction, eB> of the shock nucleon energyy density e2, i.e., BB22 = y/32TTnimpTc = K2T. (1-16) Thuss we find 4 2 uumm = nv oc r oc i£* . (i.i7) Wee refer the reader to Sari et al. (1998) and Wijers and Galama (1999) for the remaning scalings
2 i/cc ex i£ (1.18) vv&& c* *obs — constant (1.19)
FFmm oc tJto = constant. (1.20)
19 9 ChapterChapter I
Inn this simple model, the evolution of the afterglow flux, F„, at observing frequency v, can now bee determined from these four scalings (e.g., Sari et al. 1998). As both the spectrum and these scalingss are power laws, the evolution of the flux is also a power law in time. For example, for p i)/A vvmm < v < yc, the decay of the flux is F„
1.33 Afterglow
Motivatedd by this prediction of a late-time softer radiation, several groups executed rapid radio follow-upp observations of GRB error boxes (see, e.g, Chapter 2). Detection of a radio afterglow seemedd most promising. Not only does the large field of view match well with the large error boxess (several degrees) that were then available on short time scales (within a day), but maxi- mumm light was also expected to occur later at longer wavelengths. The best (pre-BeppoSAX era) limitss on such afterglow radio emission were obtained for GRB 940301. This GRB triggered an extensivee multi-wavelength campaign with ground based optical and radio observatories from thee BATSE/COMPTEL/NMSU Rapid Response Network (RRN; McNamara et al. 1995). No obviouss candidate radio counterparts were found (Frail et al. 1994; Koranyi et al. 1995; Galama ett al. 1997a, see Chapter 2).
1.3.11 The first identifications Thee breakthrough came in early 1997, when the Wide-Field Cameras (WFCs; Jager et al. 1993; seee Fig. 1.10) onboard the Italian-Dutch satellite BeppoSAX (Piro, Scarsi and Butler 1995; seee Fig. 1.11) obtained their first quickly available (within hours) accurate positions of GRBs (severall arcminutes). This allowed rapid follow-up observations which led to the discoveries off X-ray (Costa et al. 1997a), optical (Van Paradijs et al. 1997, see Chapter 5), millimeter (Bremerr et al. 1998) and radio (Frail et al. 1997a) counterparts of GRBs. This quickly settled thee distance controversy. The first transient optical counterpart, of GRB 970228, is in a faint galaxyy with ~ 0.8" diameter (Sahu et al. 1997), and the optical spectrum of GRB 970508 containss redshifted absorption lines (Metzger et al. 1997). GRBs come from 'cosmological' distances.. They are by far the most luminous photon sources in the Universe, with (isotropic) peakk luminosities in 7 rays up to 1052 erg/s, and total energy budgets up to several 1053~54 erg (Kulkarnii et al. 1998a, 1999a; see also Table 1.1). The optical signal from GRB can be 18 mag brighterr (absolute) than the brightest supernovae (Akerlof et al. 1999).
1.3.22 Confirmation of the relativistic blast-wave model Radioo light curves of the afterglow of GRB 970508 show variability on time scales of less than aa day, but these dampen out after one month (Frail et al. 1997a). Interpreting this as the effect off source expansion on the diffractive interstellar scintillation a source size of roughly 1017 cm wass derived, corresponding to a mildly relativistic expansion of the shell (Frail et al. 1997a). Thee first X-ray (Costa et al. 1997a) and optical (Galama et al. 1997b, Chapter 6, but see Galamaa et al. 1999a, Chapter 16) afterglows, discovered following GRB 970228, show a power- laww temporal decay; this is a trend observed in all subsequent X-ray and optical afterglows, with
20 0 Introduction Introduction
Figuree 1.10: BeppoSAX Wide Field Camera (Courtesy SRON) power-laww exponents in the range 1 to 2 (Groot et al. 1998a, see Chapter 12). These afterglow lightt curves agree well with the predictions of the relativistic blast-wave model (Wijers, Rees andd Mészaros 1997; Waxman 1997; Reichart 1997; see Sect. 1.2.3). Thee broad-band afterglow spectra are also power laws (in four distinct regions); together withh the observed decrease of the cooling break and the peak frequency the observations con- formm nicely with simple relativistic blast-wave models in which the emission is synchrotron radiationn by electrons accelerated in a relativistic shock (Galama et al. 1998a, see Chapter 9, Sect.. 1.2.3 and Fig. 1.9). Simplee versions of such models require as parameters the energy per unit solid angle in the blastt wave, £, the ambient density n, the electron energy spectrum (assumed to be a power law withh exponent p), and the fractions, eB and se, of energy that go into magnetic field and elec- trons,, respectively. The synchrotron afterglow spectrum of a GRB allows us to to measure five numbers:: a spectral slope (reflecting the electron energy spectrum p), three break frequencies (correspondingg to self-absorption, /^a, the peak, um, and transition to rapid electron cooling, vc), andd the flux at the peak, Fm. If the redshift, 2, is also known, all blast wave parameters can be deducedd (Wijers and Galama 1999, see Chapter 10). Thee brightness temperature of the GRB 990123 optical flash (Akerlof et al. 1999; Sect. 1.4.6)) exceeds the Compton limit of 1012 K, confirming the highly relativistic nature of the GRBB source (Galama et al. 1999b, Chapter 15).
1.44 Current status and open questions
1.4.11 GRB progenitors and the cause of the explosion Thee GRB and the afterglow are produced when relativistic ejecta are slowed down; no observ- ablee radiation emerges directly from the 'hidden engine' that powers the GRB. Thus, in spite off all recent discoveries the origin of GRBs remains mysterious. Currently popular models forr the origin of GRBs are the neutron star-neutron star (e.g., Eichler et al. 1989) and neu- tronn star-black hole mergers (Mochkovitch et al. 1993; Lattimer and Schramm 1974; Narayan,
21 1 ChapterChapter I
Paczynskii and Piran 1992), white dwarf collapse (Usov 1992), and core collapses of very mas- sivee stars ('failed' supernovae or hypernovae; Woosley 1993; Paczynski 1998). These models cann in principle provide the required energies. Observationallyy it is hard to distinguish between models of progenitors for GRBs. The locationss of GRBs in host galaxies may put constraints on such models. For instance, when aa neutron star is formed in a supernova explosion it receives a substantial kick velocity of severall hundred km/s (e.g., Van den Heuvel and van Paradijs 1998 and refs therein); together withh the relatively long merger times (~ 108 year) one would then expect some GRBs to be locatedd outside the host galaxy where the binary was formed (Bloom, Sigurdsson and Pols 1999).. Massive star collapses, on the other hand, are expected to occur in the star forming regionss where they originated. A number of GRB sources are offset from their host galaxy's centerr (Bloom et al. 1999a), but current statistics on offsets provide insufficient constraints on thee nature of GRB progenitors. Host galaxy absorption may also provide clues to the GRB progenitorr (low extinction is expected for the merger models and high extinction for massive starr collapses in star forming regions).
1.4.22 The relationship between GRBs and the star-formation rate Hostt galaxies have now been seen in most optical afterglow images (e.g, Hogg and Fruchter 1998).. The detection of [O II] A 3727 and Lyman a emission from some hosts indicates that thesee are sites of vigorous star formation (Bloom et al. 1998a; Kulkarni et al. 1998a; Djorgovski ett al. 1998c). At present the evidence for this is only suggestive; an equal number of GRB hosts havee been observed spectroscopically with Keck and no lines have been seen (private commu- nicationn Prof. S.R. Kulkarni). This might be because they are in aredshift range inaccessible to Keckk (1.4 < z < 2.4) or because they exhibit no strong star formation.
1.4.33 GRBs as potential probes of the high-redshift Universe Thee observed connection between some GRBs and star forming regions suggests that GRBs occurr at critical phases in the evolution of massive stars. Early fits to the BATSE GRB peak fluxx distribution indicated redshifts z ~ 1 for the weakest observed bursts, but these fits relied onn the assumption that the GRB rate was constant over the history of the Universe (e.g., Mao andd Paczynski 1992a). If GRBs are related to the deaths of massive stars (whose total lifetime iss very short), their rate is proportional to the star formation rate (SFR). In that case GRBs may veryy well be at very high redshifts, with z ~ 6 or greater, for the faintest bursts (e.g., Wijers ett al. 1998). A similar result was derived by Fenimore and Bloom (1995) from the observed timee dilation. GRBs may therefore become a powerful tool to probe the far reaches of the Universee by guiding us to regions of very early star formation, and the (proto) galaxies and (proto)) clusters of which they are part. The redshifts determined so far (see Table 1.1) range betweenn z — Ü.695 and z — 3.42; these correspond to peak 7-ray luminosities in the range 1051 too ÏO"3 erg/s, indicating that the GRB luminosity function is not narrow.
1.4.44 Not all GRBs have afterglows Nott all GRBs show X-ray, optical or radio afterglows (see Table 1.1). Environmental effects, suchh as absorption in a host galaxy (Groot et al. 1998b, see Chapter 11), or differences in
22 2 Introduction Introduction
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C N-'' Q. OO U SS -O (NN CS O O ww >-> ÖDD O CJ CJ 27 7 ChapterChapter 1 ON N 5 5 ON N ON N Os s ON NO N N ON N 1»ON N » ON N ON N l l 15 5 O O "5 "5 c-4 4 «« "3 X X d d NO O CM M so o NN ^ X X -H H O O O O ts s 'S Si— — e e ON N < < V V *** :> ON N en enVii c c c A A A f2 2 < CN N ^-^ ^ ijj Ê- EE — i i Os s uu 7 ^—-" " Os s 77 E Os s 4)) "4 GOO gp J22 -* «uu "5 OJ J ^^ ir- ^^ ^5 ^j B B 11 I ss Si v. v. u u OO O ^ SS S 3 3 ,, ( T—1 '—i S:: ^ s:: e "O O S S ^ ^ •—— CN c3 3 /—.. C3 a o o büü "O a >> > SP O O X X o o O O oo u ^ e e B B O O mm 00 optic optic 1-H H 28 8 Introduction Introduction Figuree 1.11: The Italian-Dutch BeppoSAX X-ray satellite. Visible are the Narrow Field Instru- mentsments (to the left) and, underneath a solar panel, one of the Wide Field Cameras. interstellarr medium densities, are plausible explanations (the brightness of the afterglow is a functionn of the ISM density; see, e.g., Sari et al. 1998; Wijers and Galama 1999). However, thee observed range in decay constants (see Table 1.1) of X-ray and optical light curves leaves roomm for very fast decays as a contributing factor (Groot et al. 1998a). Weather conditions, the sizee of the error box, and how soon an accurate position was available to observers also play an importantt role in the success of follow-up. 1.4.55 Different classes of GRBs; supernova-GRBs Thee coincidence of GRB 980425 with a peculiar type Ic supernova in a nearby galaxy raises thee possibility that very different mechanisms can give rise to GRBs (Galama et al. 1998d; Kulkarnii et al. 1998b). Modelling of the optical light curve of SN 1998bw shows that the time off collapse coincides with that of the GRB to within (+0.7,-2.0) days (Iwamoto et al. 1998). AA conservative estimate of the probability of a chance coincidence of the supernova and the GRBB is 9 x 10~5 (Galama et al. 1998d). The consequence of accepting such an association iss that the 7-ray peak luminosity of GRB 980425 and its total 7-ray energy budget are much smallerr (a factor of ~ 10°) than those of 'normal' GRBs (see Table 1.1). SN 1998bw is also peculiarr because its radio properties indicate that the radio emitter must expand relativistically (Kulkarnii et al. 1998b; Li and Chevalier 1999; but see Waxman and Loeb 1999). But, perhaps most,, if not all, GRBs are associated with supernovae; Bloom et al. (1999b) attribute the late- timee light curve of GRB 980326 to a supernova. Also for GRB 970228 there is evidence that aa supernova dominated the light curves at late times (Reichart 1999; Galama et al. 1999a, see Chapterr 16). So there may not be a dichotomy between 'normal' and supernova GRBs, only a graduall transition. 29 9 ChapterChapter 1 1.4.66 The early afterglow Thee discovery of a very bright and brief optical flash coincident in time with GRB 990123 (Akerloff et al. 1999) shows that the early optical signal from GRB can be some 18 magnitudes brighterr than the brightest supernovae. It was predicted that the reverse shock could cause emissionn (Mészaros and Rees 1997; Sari and Piran 1999a) that peaks in the optical waveband andd is observed only during or just after the GRB. GRB 990123 would then be the first burst in whichh all three emitting regions have been seen: internal shocks causing the GRB, the reverse shockk causing the prompt optical flash, and the forward shock causing the afterglow (Mészaros andd Rees 1999; Sari and Piran 1999b; Galama et al. 1999b, see Chapter 15). The brief radio flareflare observed for GRB 990123 (Kulkarni et al. 1999b; Galama et al. 1999b) has been suggested too be due to the reverse shock (Sari and Piran 1999b; Kulkarni et al. 1999b). However, an interpretationn in terms of emission by the forward shock (Galama et al. 1999b) is also consistent withh the observations. 1.4.77 Strongly anisotropic outflow (beaming) Ann important uncertainty concerns the possible beaming of the 7-ray and afterglow emissions. Thiss has an immediate impact on the burst energetics, the nature and number of events needed too account for the observed burst rate, and the prospects for optical detection of afterglows fromm events whose 7-ray beam does not point toward us (Rhoads 1997, 1999). If the afterglow iss beamed with opening angle 9, an extra change of the light curve slope occurs at the time whenn the Lorentz factor T of the blast wave equals 1/9. Slightly later the jet begins a lateral expansion,, which causes a further steepening of the light curve. Perhaps such a transition hass been observed in the optical afterglow light curve of GRB 990123 (Kulkarni et al. 1999a; Fruchterr et al. 1999). A similar transition was better sampled in afterglow data of GRB 990510; Harrisonn et al. (1999) and Stanek et al. (1999) present optical observations of GRB 990510, whichh show a clear steepening of the rate of decay of the light between ~ 3 hours and several days.. Together with radio observations (Harrison et al. 1999), which reveal a similar steepening off the decline, it is found that the transition is very much frequency-independent; this virtually excludess explanations in terms of the passage of the cooling or the peak frequency, but is what iss expected in case of beaming. Harrison et al. (1999) derive a jet opening angle of 9 = 0.08, whichh for this burst would reduce the total energy in 7 rays to ~ 1051 erg. 1.55 Overview of this thesis Thiss thesis presents the search for GRB counterparts (Chapters 2 - 4) and the discovery of the firstfirst optical counterpart to a GRB (Chapter 5). After the successful identifications of X-ray, opticall and radio afterglows the research described in this thesis has concentrated on obtaining andd interpreting multi-wavelength observational information on GRB afterglows (Chapters 5 - 16).. I also present the discovery of a possible new class of GRB sources which are associated withh supernovae (S-GRBs; Chapter 13). A summary in Dutch (Chapter 17) and a list of my publicationss are given at the end of the thesis. Inn this Chapter (Chapter 1) I have given a general introduction to GRBs, presented some observationall characteristics of GRBs, discussed the distance scale controversy and the predic- tionn that GRBs should be acompanied by a late-time afterglow. I also discussed the discovery 30 0 Introduction Introduction off GRB afterglows, the implications of this discovery and the current status of GRB afterglow research. . Motivatedd by the prediction of low-energy afterglow (Sect. 1.2.3), we executed rapid radio follow-upp observations of GRB error boxes. Chapter 2 presents the results of a search for a radio counterpartt in the COMPTEL error box region of GRB 940301, using the Westerbork Synthesis Radioo Telescope (WSRT). We detected no flux density variations with amplitudes greater than -F325MHZZ = 10 mJy (5 a) within a period of 1 to 4 months after the burst. Inn Chapter 3 we report on a serendipitous discovery. As it turned out, the GRB 940301 fieldd contains the radio pulsar PSR B0655+64 for which we discovered a large amplification of thee 325 MHz flux density (a factor of ~ 43). The phenomenon is unrelated to GRB 940301 ass the pulsar is outside the IPN localization for this burst. The decorrelation bandwidth (~ 5 MHz)) and characteristic time scale of the variation (~ 1 hour) are in agreement with the values expectedd for diffractive interstellar scintillation. To the best of our knowledge, this is the largest suchh increase ever found for a radio pulsar. Such an extreme flux density amplification cannot bee explained by ordinary scintillation, but might be in terms of a caustic. Inn Chapter 4 we present the search for a counterpart in the BeppoSAX WFC error box of thee first rapidly localized GRB (GRB 970111) using the WSRT. We then decided to also include opticall observations in our efforts. Like other groups, we did not find radio or optical afterglow forr this burst. Chapterr 5 presents the discovery of the first optical counterpart to a GRB (GRB 970228). Thee optical transient appears to be associated with a faint galaxy. This supports a cosmological originn for this GRB. Inn Chapter 6 we construct an optical light curve of the optical counterpart of GRB 970228. Wee find that between 21 hours and six days after the burst, the R-band brightness decreased by aa factor ~ 40, with any subsequent decrease decrease in brightness occuring at a much slower rate.. As the point source faded, it also became redder. The initial behaviour of the source appearss to be consistent with the 'fireball' model, but the subsequent decrease in the rate of fadingg may prove harder to explain. Chapterss 7 and 8 present observations of GRB 970508 with the William Herschel Telescope (WHT;; U, B, V, ^ and Ic; Chapter 7) and with the WSRT (1.4 GHz; Chapter 8). The optical lightt curve peaks around 1.9 days, after which it follows a decline that can be fitted well with aa single power law. We show that deviations from a smooth power-law decay are only mod- eratee (r.m.s. =0.15 magnitude). We find that the optical spectral energy distribution changed aroundd the time of the maximum (the spectrum became redder). Our WSRT observations show aa transition from optically thick to optically thin emission ~ 50 days after the event. Inn Chapter 9 we show evidence for synchrotron emission, and the existence and importance off the cooling frequency vc in the afterglow of GRB 970508. We present the reconstructed broad-bandd GRB 970508 afterglow spectrum 12 days after the event. The spectrum beautifully supportss the idea that the afterglow is synchrotron emission by relativistic electrons. Sari et al.. (1998) and Mészaros, Rees and Wijers (1998) argue that, apart from the synchroton peak frequencyy um, the cooling frequency vc\s important for the temporal and spectral evolution of thee afterglow. We realized that the observed reddening of the optical light and the observed relationn between the power-law spectral and temporal slope is in agreement with such a model (seee Chapter 8). In our model the observed spectral transition in the optical and infrared is explainedd by the cooling frequency vc passing through these passbands; maximum light at 86 GHzz (Bremer et al. 1998) and at 8.5 GHz (Frail et al. 1997a) is explained by the passage of the 31 1 ChapterChapter 1 peakk frequency um. Inn Chapter 10 we use the properties, deduced from the light curve behavior (Chapters 7 and 8)) and broad-band spectrum (Chapter 9) of GRB 970508, to calculate the physical parameters off the afterglow. Dr. Wijers and I realized that, for the first time, all model parameters could bee deduced for GRB 970508 with these results. We calculate synchrotron spectra of relativistic blastt waves and give expressions to infer the physical properties of the afterglow from the measuredd spectral features. We compute the ambient density, n = 0.03 cm"3, and the blast- wavee energy per unit solid angle, € = 3 x 1052erg/47r sr. We also computed the energy density inn electrons and magnetic field, and find that they are 12% and 9%, respectively, of the nucleon energyy density. Inn Chapter 11 we report the absence of any transient optical source brighter than Rc = 23.8 magg in observations of the error box of GRB 970828, and discuss the consequences of this non- detectionn and the possible effect of redshift on the relation between optical absorption and the low-energyy cut off in the X-ray afterglow spectrum. Chapterr 12 describes our discovery of the optical counterpart to GRB 980326, which dis- playedd the most rapid optical decay observed so far (a power law with exponent 2.10 0.13). Thee rapid decay may be a reason for the non-detection of some low-energy afterglows of GRBs. Inn Chapter 13 we report the discovery of a luminous peculiar type Ic supernova, SN 1998bw, inn the error box of GRB 980425, which occurred within about a day of the GRB. The light curvee is very different from those of known optical afterglows of GRBs. The consequence off an association is that the 7-ray peak luminosity of GRB 980425 and its total 7-ray energy budgett are much smaller (a factor of ~ 105) than those of 'normal' GRBs (see Table 1.1). Our discoveryy raises the possibility that very different mechanisms may give rise to GRBs, which differr little in their 7-ray properties. Inn Chapter 14 we present the discovery of the X-ray afterglow of GRB 980703 and report onn X-ray, optical and infrared follow-up observations of GRB 980703. We find no evidence forr a spectral break in the infrared to X-ray spectral range on 1998 July 4.4, and determine a 17 lowerr limit of the cooling break frequency: vc > 1.3 x 10 Hz. Comparison of the afterglow off GRB 980703 with that of GRB 970508 yields that the fraction of the energy in the magnetic 5 field,field, tB < 6 x 10 , is much lower for this burst. Inn Chapter 15 we report on the results of 7-ray, optical, infrared, submillimetre and radio observationss of GRB 990123 and its afterglow. The discovery of a very bright and brief opical flashh coincident in time with this GRB (Akerlof et al. 1999) allows a comparison of this early flashflash with the GRB and with the late-time multiwavelength observations of its afterglow. Our interpretationn of the data indicates that the initial and afterglow emissions are associated with threee distinct regions in the fireball. We suggest that the differences between bursts reflect variationss in the magnetic-field strengths in the afterglow-emitting regions. Inn Chapter 16 we give an update and reevaluation of the light curve of the optical transient, associatedd with GRB 970228. We confirm the result of Galama et al. (1997b; Chapter 6) that thee early decay of the light curves (before March 6, 1997) was faster than that at later times (betweenn March 6 and April 7, 1997). The early-time observations of GRB 970228 are consis- tentt with relativistic blast-wave models but the late-time observations are hard to understand in thiss framework. The observations are well explained by an initial power-law decay with a = —— 1.73+0 12 modified at later times by a type-Ic supernova light curve. This is further evidence thatt at least some GRBs are associated with supernovae. 32 2 Introduction Introduction Acknowledgments.Acknowledgments. 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