GRB 061121: Broadband spectral evolution through the prompt and
  afterglow phases of a bright burst
Draft version November 21, 2018
Preprint typeset using LATEX style emulateapj v. 08/22/09
GRB 061121: BROADBAND SPECTRAL EVOLUTION THROUGH THE PROMPT AND AFTERGLOW
PHASES OF A BRIGHT BURST.
K.L. Page1, R. Willingale1, J.P. Osborne1, B. Zhang2, O. Godet1, F.E. Marshall3, A. Melandri4, J.P. Norris5,6,
P.T. O’Brien1, V. Pal’shin7, E. Rol1, P. Romano8,9, R.L.C. Starling1, P. Schady10, S.A. Yost11, S.D. Barthelmy3,
A.P. Beardmore1, G. Cusumano12, D.N. Burrows13, M. De Pasquale10, M. Ehle14, P.A. Evans1, N. Gehrels3, M.R.
Goad1, S. Golenetskii7, C. Guidorzi8,9, C. Mundell4, M.J. Page10, G. Ricker15, T. Sakamoto3, B.E. Schaefer16,
M. Stamatikos3, E. Troja1,12, M.Ulanov7, F. Yuan11 & H. Ziaeepour9
Draft version November 21, 2018
ABSTRACT
Swift triggered on a precursor to the main burst of GRB 061121 (z = 1.314), allowing observations to be made
from the optical to gamma-ray bands. Many other telescopes, including Konus-Wind, XMM-Newton, ROTSE and
the Faulkes Telescope North, also observed the burst. The gamma-ray, X-ray and UV/optical emission all showed a
peak ∼ 75 s after the trigger, although the optical and X-ray afterglow components also appear early on – before,
or during, the main peak. Spectral evolution was seen throughout the burst, with the prompt emission showing a
clear positive correlation between brightness and hardness. The Spectral Energy Distribution (SED) of the prompt
emission, stretching from 1 eV up to 1 MeV, is very flat, with a peak in the flux density at ∼ 1 keV. The optical-
to-X-ray spectra at this time are better fitted by a broken, rather than single, power-law, similar to previous results
for X-ray flares. The SED shows spectral hardening as the afterglow evolves with time. This behaviour might be
a symptom of self-Comptonisation, although circumstellar densities similar to those found in the cores of molecular
clouds would be required. The afterglow also decays too slowly to be accounted for by the standard models. Although
the precursor and main emission show different spectral lags, both are consistent with the lag-luminosity correlation
for long bursts. GRB 061121 is the instantaneously brightest long burst yet detected by Swift. Using a combination of
Swift and Konus-Wind data, we estimate an isotropic energy of 2.8 × 1053 erg over 1 keV – 10 MeV in the GRB rest
frame. A probable jet break is detected at ∼ 2 × 105 s, leading to an estimate of ∼ 1051 erg for the beaming-corrected
gamma-ray energy.
Subject headings: gamma-rays: bursts — X-rays: individual (GRB 061121)
1. INTRODUCTION
Electronic address: kpa@star.le.ac.uk
1 Department of Physics and Astronomy, University of Leicester,
Leicester, LE1 7RH, UK
2 Department of Physics & Astronomy, University of Nevada,
Las Vegas, NV 89154-4002, USA
3 NASA/Goddard Space Flight Center, Greenbelt, MD 20771,
4 Astrophysics Research Institute, Liverpool John Moores Uni-
versity, Twelve Quays House, Egerton Wharf, Birkenhead, CH41
5 Denver Research Institute, University of Denver, Denver, CO
80208, USA
6 Visiting Scholar, Stanford University
7 Ioffe Physico-Technical Institute, Laboratory for Experimen-
tal Astrophysics, 26 Polytekhnicheskaya, Saint Petersburg 194021,
Russian Federation
8 INAF, Osservatorio Astronomico di Brera, Via E. Bianchi 46,
I-23807, Merate (LC), Italy
9 Dipartimento di Fisica, Universitá di Milano-Bicocca, Piazza
delle Scienze 3, I-20126, Milano, Italy
10 Mullard Space Science Laboratory, University College Lon-
don, Holmbury St. Mary, Dorking, Surrey RH5 6NT, UK
11 University of Michigan, 2477 Randall Laboratory, 450 Church
St., Ann Arbor, MI 48104, USA
12 INAF-IASF, Sezione di Palermo, via Ugo La Malfa 153, 90146,
Palermo, Italy
13 Department of Astronomy and Astrophysics, Pennsylvania
State University, 525 Davey Lab, University Park, PA 16802, USA
14 XMM-Newton Science Operations Centre, European Space
Agency, Villafranca del Castillo, Apartado 50727, E-28080 Madrid,
Spain
15 Center for Space Research, Massachusetts Institute of Tech-
nology, 70 Vassar Street, Cambridge, MA 02139, USA
16 Department of Physics and Astronomy, Louisiana State Uni-
versity, Baton Rouge, LA 70803, USA
Gamma-Ray Bursts (GRBs) are intrinsically extremely
luminous objects, approaching values of 1054 erg s−1 if
the radiation is isotropic (e.g., Frail et al. 2001; Bloom
et al. 2003). This energy is emitted over all bands in the
electromagnetic spectrum; to understand GRBs as fully
as possible, panchromatic observations are required over
all time frames of the burst.
The Swift multi-wavelength observatory (Gehrels et
al. 2004) is designed to detect and follow-up GRBs.
With its rapid slewing ability, Swift is able to follow
bursts and their afterglows from less than a minute af-
ter the initial trigger, and can often still detect them
weeks, and sometimes months, later. On rare occa-
sions, such as when Swift triggers on a precursor to the
main burst, the prompt emission, as well as the after-
glow, can be observed at X-ray and UV/optical wave-
lengths. GRB 061121, the subject of this paper, is only
the third GRB Swift has detected in this manner (af-
ter GRB 050117 – Hill et al. 2006 and GRB 060124 –
Romano et al. 2006), out of the almost 200 bursts trig-
gered on in the first two years of the mission.17 Of
these, GRB 061121 is the second well-sampled event
(GRB 060124 was the first), and the first for which the
UV/Optical Telescope (UVOT) was in event mode.
In addition to the small number of precursor triggers,
around 10% of Swift bursts show detectable emission over
17 GRB 050820A would possibly have also been in this category,
but Swift entered the South Atlantic Anomaly (SAA) just as a
dramatic increase in count rate began (Cenko et al. 2006; Page et
al. 2005a; Cummings et al. 2005; Page et al. 2005b; Chester et al.
2005); Swift does not actively collect data during SAA passages.
http://arxiv.org/abs/0704.1609v1
mailto:kpa@star.le.ac.uk
2 K.L. Page et al.
the BAT bandpass by the time the narrow field instru-
ments (NFIs) are on target.
Besides the Swift observations of prompt emission,
there have been a small number of prompt optical mea-
surements of GRBs, thanks to the increasing number
of robotic telescopes around the world. A variety of
behaviours has been found, with some optical (and in-
frared) light-curves tracking the gamma-ray emission
(e.g., GRB 041219A – Vestrand et al. 2005; Blake
et al. 2005), while others appear uncorrelated (e.g.,
GRB 990123 – Akerlof et al. 1999, Panaitescu & Kumar
2007, though see also Tang & Zhang 2006; GRB 050904
– Boër et al. 2006; GRB 060111B – Klotz et al. 2006;
GRB 060124 – Romano et al. 2006). GRB 050820A (Ves-
trand et al. 2006) showed a mixture of both correlated
and uncorrelated optical flux.
Where correlations exist between different energy
bands, it is likely that there is a common origin for
the components. In the uncorrelated cases, the opti-
cal emission may be due to an external reverse shock
(e.g., Sari & Piran 1999; Mészáros & Rees 1999), while
the prompt gamma-rays are caused by internal shocks.
Cenko et al. (2006) suggest that the early optical data
for GRB 050820A are produced by the forward shock
passing through the band. In the case of GRB 990123,
Panaitescu & Kumar (2007) have suggested that the
gamma-rays arose from inverse Comptonisation, while
the optical emission was due to synchrotron processes;
they do not assume a specific mechanism for the energy
dissipation, allowing for the possibility of either internal
or reverse-external shocks.
It is unclear whether precursors are ubiquitous fea-
tures of GRBs, often remaining undetected because of
a low signal-to-noise ratio or being outside the energy
bandpass of the detector, or whether only some bursts
exhibit them. A detailed discussion of the precursor phe-
nomenon is beyond the scope of this paper and will be
addressed in a future publication.
In this paper, we report on the multi-wavelength ob-
servations of both the prompt and afterglow emission of
GRB 061121. §2 details the observations made by Swift,
Konus-Wind, XMM-Newton, ROTSE18 and the Faulkes
Telescope North (FTN), with multi-band comparisons
being made. In §3, we discuss the precursor, prompt and
afterglow emission, with a summary given in §4.
Throughout the paper, the main burst (∼ 60–200 s
after the trigger) will be referred to as the prompt emis-
sion, and the emission seen over −5 to +10 s as the
precursor, where the BAT trigger time T0 = 0 s. Er-
rors are given at 90% confidence (e.g., ∆χ2 = 2.7 for
one interesting parameter) unless otherwise stated, and
the convention Fν,t ∝ ν
−βt−α (with the photon spec-
tral index, Γ = β + 1 where dN/dE ∝ E−Γ) has been
followed. We have assumed a flat Universe, with Hubble
constant, H0 = 70 km s
−1 Mpc−1, cosmological constant,
ΩΛ = 0.73 and Ωmatter = 1−ΩΛ.
2. OBSERVATIONS AND ANALYSES
Two years and one day after launch, the Burst Alert
Telescope (BAT; Barthelmy et al. 2005) triggered on a
precursor to GRB 061121 at 15:22:29 UT on 21st Novem-
ber, 2006. Swift slewed immediately, resulting in the
18 Robotic Optical Transient Search Experiment
NFIs being on target and beginning to collect data 55 s
(X-ray Telescope: XRT; Burrows et al. 2005a) and 62 s
(UVOT; Roming et al. 2005) later. This enabled broad-
band observations of the main burst event, which peaked
∼ 75 s after the trigger, leading to spectacular multi-
wavelength coverage of the prompt emission. The most
accurate Swift position for this burst was that determined
by the UVOT: RA = 09h 48m 54.s55, decl = −13◦ 11′ 42.′′4
(J2000.0; 90% confidence radius of 0.′′6; Marshall et al.
2006); the refined XRT position is only 0.′′1 from these
coordinates (Page et al. 2006b).
GRB 061121 was declared a ‘burst of interest’ by the
Swift team (Gehrels et al. 2006a), to encourage an inten-
sive ground- and space-based follow-up programme. In
addition to the Swift observations, the prompt emission
of GRB 061121 was detected by RHESSI19 (Bellm et
al. 2006), Konus-Wind and Konus-A (Golenetskii et al.
2006). Later afterglow observations were obtained in the
X-ray (XMM-Newton – Schartel 2006) and radio (VLA20
– Chandra & Frail 2006) bands. ATCA21 and WSRT22
also observed in the radio band between ∼5.2 day and
∼6.2 day after the burst, but did not detect the after-
glow (van der Horst et al. 2006a,b), implying it had faded
since the VLA observation.
Likewise, extensive optical follow-up observations were
performed: ROTSE-IIIa (Yost et al. 2006), FTN (Me-
landri et al. 2006), Kanata 1.5-m telescope (Uemura et
al. 2006), the University of Miyazaki 30-cm telescope
(Sonoda et al. 2006), MDM23 (Halpern et al. 2006a,b;
Halpern & Armstrong 2006a,b), P6024 (Cenko 2006),
ART25 (Torii 2006), the CrAO26 2.6-m telescope (Efimov
et al. 2006a,b) and SMARTS/ANDICAM27 (at infrared
wavelengths, too; Cobb 2006) all detected the optical
afterglow. Spectroscopic observations were performed
at the Keck telescope about 12 minute after the trig-
ger, finding a redshift of z = 1.314 for the optical af-
terglow, based on absorption features (Perley & Bloom
2006; Bloom et al. 2006).
GRB 061121 has the highest instantaneous peak flux
of all the long bursts detected by Swift to date (e.g.,
Angelini et al. in prep).
2.1. Gamma-ray Data
2.1.1. BAT
Temporal Analysis— After the initial precursor, the BAT
count rate returned to close to the instrumental back-
ground level, until T0+60 s, at which point the much
brighter main burst began. This is characterised by a
series of overlapping peaks, each brighter than the previ-
ous one, after which the gamma-ray flux decayed (from
∼T0+75 s to ∼T0+140 s). Event data were collected un-
til almost 1 ks after the trigger, thus covering the entire
emission period.
19 Reuven Ramaty High Energy Solar Spectroscopic Imager
20 Very Large Array
21 Australia Telescope Compact Array
22 Westerbork Synthesis Radio Telescope
23 Michigan-Dartmouth-MIT Observatory
24 Palomar 60 inch
25 Automated Response Telescope
26 Crimean Astrophysical Observatory
27 Small and Moderate Aperture Research Telescope System/A
Novel Double-Imaging CAMera
GRB 061121: Broadband observations 3
T90, over 15-150 keV, and incorporating both the pre-
cursor and main emission, is 81 ± 5 s, measured from
8.8–89.8 s after the trigger28. Figure 1 shows the mask-
weighted BAT light-curve in the four standard energy
bands [15–25, 25–50, 50–100, 100-150 keV; 64 ms bin-
ning between 50-80 s after the trigger, with 1 s bins
at all other times; units of count s−1 (fully illuminated
detector)−1], with light-curves from other instruments:
the precursor and the pulses of the main burst are de-
tected over all gamma-ray bands, although the precursor
is only marginal over the 100-150 keV BAT band. There
is also a soft tail (detected below ∼ 50 keV, when suffi-
ciently coarse time bins are used) visible until about 140 s
after the trigger (see bottom panel of Figure 1), corre-
sponding to a similar feature in the X-ray light-curves.
Spectral analysis— For the precursor,
T90,pre = 7.7 ± 0.5 s (15–150 keV). A spectrum ex-
tracted over this interval can be well modelled by a single
power-law, with Γ = 1.68 ± 0.09 (χ2/dof = 26.2/23); no
significant improvement was found by using the Band
function (Band et al. 1993) or a cut-off power-law and a
thermal model led to a slightly (χ2 ∼ 8) worse fit. The
15–150 keV fluence for this time interval is 4 × 10−7
erg cm−2.
Considering only the main event,
T90,main = 18.2 ± 1.1 s (measured from 61.8–80.0 s
post-trigger). Fitting a power-law to the mean
spectrum during this time also results in a good fit
(Γ = 1.40 ± 0.01; fluence = 1.31 × 10−5 erg cm−2 over
15–150 keV; χ2/dof =51.6/56 ); again, neither the Band
function nor a cut-off power-law improves upon this.
There is significant spectral evolution during the T90
period, as shown in Figure 2: at times when the count
rate is higher, the spectrum is harder. This behaviour
was also common in earlier bursts, as well as previous
Swift detections (e.g. Golenetskii et al. 1983; Ford et al.
1995; Borgonovo & Ryde 2001; Goad et al. 2007). The
precursor shows a similar dependence of hardness ratio
on count rate, suggesting that the emission processes
in the precursor and the main burst are the same or
similar.
2.1.2. Konus-Wind
Temporal Analysis— Konus-Wind (Aptekar et al. 1995)
triggered on the main episode of GRB 061121, while
Konus-A triggered on the precursor (Golenetskii et
al. 2006). Because of the spatial separation of
Swift and Wind, the light travel-time between the
spacecraft is 1.562 s: the Konus-Wind trigger time,
T0,K−W = T0,BAT + 61.876 s. All Konus light-curves
have been plotted with respect to the BAT trigger,
corrected for the light travel-time. Figure 1 shows
the Konus-Wind data plotted over the standard energy
bands, with 64 ms binning; the bottom panel plots the
coarser time resolution (2.944 s) ‘waiting mode’ data,
showing that Konus-Wind did see slightly enhanced
emission at the time of the precursor. The background
levels (which have been subtracted in each case) were
1005, 370 and 193.4 count s−1 for bands 21–83, 83–360
and 360–1360 keV, respectively.
28 Errors on the BAT T90 are estimated to be typically 5–10%,
depending on the shape of the light-curve.
Spectral analysis— Table 1 gives the spectral fits to
the Konus-Wind data in three separate time intervals
shown by vertical lines in Figure 1 (Konus-Wind spec-
tral intervals are automatically selected onboard): up
to the end of the ‘bump’ around 70 s (the ‘start’
of the burst), the burst maximum and, finally, until
most of the emission has died away (the burst tail).
The data were fitted with a cut-off power-law, where
dN/dE ∼ E−Γ × e[−(2−Γ)E/Epeak], leading to the photon
indices and peak energies given in the table. The Band
function was used to estimate upper limits for the pho-
ton index above the peak; the values for the peak energy
and Γ obtained from the Band function were the same
as when fitting the cut-off power-law. Little variation in
the spectral slope for energies below the peak is seen over
these intervals, though the peak itself may have moved to
somewhat higher energies during the burst emission. Ex-
tracting BAT spectra over the same time intervals, and
fitting with the same model (fixing Epeak at the value
determined from the Konus-Wind data) results in con-
sistent spectral indices.
2.2. X-ray Data
2.2.1. XRT
Temporal Analysis— The XRT identified and centroided
on an uncatalogued X-ray source in a 2.5 s Image Mode
(IM) frame, as soon as the instrument was on target.
This was quickly followed by a pseudo Piled-up Photo
Diode (PuPD) mode frame. Following damage from a
micrometeoroid impact in May 2005 (Abbey et al. 2005),
the Photo Diode mode (Low Rate and Piled-up) has been
disabled [see Hill et al. (2004) for details on the different
XRT modes]; however, the XRT team are currently work-
ing on a method to re-implement these science modes and
to update the ground software to process the files. The
pseudo PuPD point presented here is the first use of such
data.
Data were then collected in Windowed Timing (WT)
mode starting at a count rate of ∼ 1280 count s−1 (pile-
up corrected – see below); the rate rapidly increased to
a maximum of ∼ 2500 count s−1 at T0 + 75 s, mak-
ing GRB 061121 the brightest burst yet detected by the
XRT. Following this peak, the count-rate decreased, with
a number of small flares superimposed on the underlying
decay (see Figure 1). Photon Counting (PC) mode was
automatically selected when the count rate was below
about 10 count s−1. Around 1.5 ks, the XRT switched
back into WT mode briefly, due to an enhanced back-
ground linked to the sunlit Earth and a relatively high
CCD temperature.
Because of the high count rate, the early WT data were
heavily piled-up; see Romano et al. (2006) for informa-
tion about pile-up in this mode. To account for this, an
extraction region was used which excluded the central
20 pixels (diameter; 1 pixel = 2.′′36) and extended out
to a total width of 60 pixels. Likewise, the first three
orbits of PC data were piled-up, and the data were thus
extracted using annular regions (inner exclusion diame-
ter decreasing from 12 to 6 to 4 pixels as the afterglow
faded; outer diameter 60 pixels). The count rate was
then corrected for the excluded photons by a comparison
of the Ancillary Response Files (ARFs) generated with
and without a correction for the Point Spread Function
4 K.L. Page et al.
White UVOT
0.3−2 keV XRT
2−10 keV
15−25 keV BAT
25−50 keV
50−100 keV
100−150 keV
21−83 keV Konus
83−360 keV
360−1360 keV
0 50 100 150
time since BAT trigger (s)
21−1360 keV
(waiting mode)
50 100 150
time since BAT trigger (s)
15−50 keV
Fig. 1.— Top panels: Swift UVOT, XRT, BAT and Konus-Wind light-curves of GRB 061121; 1σ error bars are shown for the UVOT
and XRT data. Each instrument detected the peak of the main burst, with the precursor being detected over all gamma-ray energies. The
vertical lines in the 360–1360 keV panel indicate the start and stop times for the spectra given in Table 1. Bottom panel: The 15-50 keV
BAT light-curve, with 10-s bins, showing a tail out to ∼140 s.
GRB 061121: Broadband observations 5
start time (s) stop time (s) Γ Epeak (keV) ΓBand χ
2/dof
61.876 70.324 1.40
+0.08
−0.09
<2.1 72/75
70.324 75.188 1.23
+0.05
−0.06
<2.9 88/75
75.188 83.380 1.30
+0.11
−0.13
<2.3 81/75
61.876 83.380 1.32
+0.04
−0.05
<2.7 95/75
TABLE 1
Konus-Wind cut-off power-law spectral fit results. Times are given with respect to the BAT trigger. ΓBand is the upper
limit obtained for the spectral index above Epeak when fitting with the Band function.
6 K.L. Page et al.
65 70 75 80
time since trigger (s)
0 1 2 3
BAT count rate (count s−1 detector−1)
Fig. 2.— Top panels: Light-curves, hardness ratios (HR) and the
variation in Γ using a single power-law fit during the main emission.
The BAT light-curve (top panel) is in units of count s−1 (fully il-
luminated detector)−1 , and the corresponding hardness ratio plots
(50–150 keV)/(15–50 keV) using 1-s binning. The XRT light-curve
shows counts over 0.3–10 keV, while the hardness ratio compares
(1–10 keV)/(0.3–1 keV) over 1-s bins. Bottom panel: BAT hard-
ness ratio versus count rate, showing that the emission is harder
when brighter. Data from the precursor are shown as grey circles,
with the main burst in black. The grey line shows a fit to the data,
of the form HR = 0.14 CR + 0.39.
(PSF); the ratio of these files provides an estimate of the
correction factor. Nousek et al. (2006) give more details
on this method. Occasionally, the afterglow was partially
positioned over the CCD columns disabled by microm-
eteoroid damage mentioned above. In these cases, the
data were corrected using an exposure map.
From T0 + 3 × 10
5 s onwards, the afterglow had faded
sufficiently for a nearby (41.′′5 away), constant (count
rate ∼ 0.003 count s−1) source to contaminate the GRB
region; this source is coincident with a faint object in
the Digitized Sky Survey and is marginally detected in
the UVOT V filter. Thus, beyond this time, the ex-
traction region was decreased to a diameter of 30 pix-
els, and the count rates corrected for the loss in PSF (a
factor of ∼ 1.08). The spectrum of this nearby source
can be modelled with a single power-law of Γ = 1.5+0.2
−0.1,
100 1000 104 105 1061
time since trigger (s)
Fig. 3.— Swift-XRT light-curve of GRB 061121. The star and
triangle show the initial Image Mode and pseudo PuPD point (see
text for details), followed by WT mode data (black) during the
main burst (and at the end of the first orbit) and PC mode data
(in grey).
10050 200 500
time since trigger (s)
Fig. 4.— Swift flux light-curve of GRB 061121, showing the early
X-ray data (star, triangle and crosses) and the BAT data (grey
histogram) extrapolated into the 0.3–10 keV band pass in units
of erg cm−2 s−1, together with the UVOT flux density light-curve
(light grey circles – V -band; dark grey circles – White filter) in units
of erg cm−2 s−1 Å−1, scaled to match the XRT flux observed at
the start of the ‘plateau’ phase.
with NH = (1.8
−1.2) × 10
21 cm−2, in comparison with
the Galactic value in this direction of 5.09 × 1020 cm−2
(Dickey & Lockman 1990).
Figure 3 shows the XRT light-curve, starting with the
IM point (see Hill et al. 2006 for details on how IM data
are converted to a count rate) and followed by the pseudo
PuPDmode data. The importance of these early pre-WT
data is clear, confirming that the XRT caught the rise of
the main burst.
After the bright burst, the afterglow began to follow
the ‘canonical’ decay, seen in many Swift bursts (Nousek
et al. 2006; Zhang et al. 2006a). Such a decay can
be parameterised by a series of power-law segments; in
this case, fitting the data beyond 200 s after the trig-
ger (= 125 s after the main peak), two breaks in the
light-curve were identified, with the decay starting off
very flat (α = 0.38 ± 0.08) and eventually steepening
GRB 061121: Broadband observations 7
α1 0.38 ± 0.08 Plateau phase
Tbreak,1 2258
α2 1.07
+0.04
−0.06
Shallow phase
Tbreak,2 (3.2
) × 104 s
α3 1.53
+0.09
−0.04
Steep phase
TABLE 2
XRT power-law light-curve fits from 200 s after the
trigger onwards; times are referenced to the BAT
trigger. The names used in the text for the different
epochs of the light-curve are listed in the last column.
to α = 1.07+0.04
−0.06 at ∼ 2.3 ks and then α =1.53
+0.09
−0.04 at
∼ 32 ks (Table 2). The addition of the second break
vastly improved the fit by ∆χ2 = 112.4 for two degrees
of freedom. However, we note that O’Brien et al. (2006)
and Willingale et al. (2007) advocate a different descrip-
tion of the temporal decline; we return to this in §3.
Fitting the decay of the main peak (75–200 s, keeping
T0 as the trigger time) with a power-law, the slope is
very steep, with α0 = 5.1 ± 0.2. However, both Zhang
et al. (2006a) and Liang et al. (2006) have shown that
the appropriate time origin is the start of the last pulse.
Thus, a model of the form f(t) ∝ (t−t0)
−α0 was used,
finding t0 = 58 ± 1 s and a slope of α0 = 2.2
−0.3; this
is a statistically significant improvement on the power-
law fit using the precursor T0 (∆χ
2 = 32 for one extra
parameter).
Figure 4 plots the Swift data in terms of flux (the BAT
data have been extrapolated into the 0.3–10 keV band,
using the joint fits with the XRT described in §2.4.1) and
flux density for UVOT. The BAT and XRT data are fully
consistent with each other at all overlapping times.
Spectral Analysis— The XRT data also show that strong
spectral evolution was present throughout the period of
the prompt emission; this is discussed in conjunction
with the BAT data in §2.4.1. Considering the X-ray data
alone, there is some indication that the spectra may be
better modelled with a broken, rather than single, power-
law, although the break energies cannot always be well
constrained (see Figure 5). For each spectrum [covering
periods of 2 s during the main pulse, followed by two
spectra of 5 s (80–85 s and 85–90 s) where the emission
is fainter], the low-energy slopes were tied together for
each spectrum (i.e., the slope measured is that averaged
over all of the spectra), as were the high-energy indices,
and the rest-frame column density, NH,z, was fixed at
(9.2 ± 1.2) × 1021 cm−2 from the best fit to the data
from later times (see below); only the break energy and
the normalisation were allowed to vary. When simul-
taneously fitting all 11 spectra, χ2/dof decreased from
142/134 to 127/132. Individually, the spectral fits were
typically improved by χ2 of between 2–5.
The X-ray data during the GRB 051117A flares (Goad
et al. 2007) were found to be better modelled with bro-
ken power-laws, with the break energy moving to harder
energies during each flare rise, and then softening again
as the flux decayed. Likewise, Guetta et al. (2006) found
breaks in the X-ray spectra obtained during the flares in
GRB 050713A. The same pattern may be occurring here,
and there is certainly an indication of spectral curvature.
The observed flux calculated from the spectrum corre-
sponding to the peak of the emission (74–76 s) was mea-
70 80 90
time since trigger (s)
70 80 90
time since trigger (s)
70 80 90
time since trigger (s)
Fig. 5.— Fitting the X-ray data over 0.3–10 keV with a broken
power-law (Γ1 =0.69
+0.13
−0.07
and Γ2 =1.61
+0.14
−0.13
for all spectra), the
break energy seems to move through the band, towards higher
energies when the emission is brighter. Arrows indicate upper or
lower 90% limits.
sured to be 1.66 × 10−7 erg cm−2 s−1 (over 0.3–10 keV);
the unabsorbed value was 1.77 × 10−7 erg cm−2 s−1.
The PC spectra were also extracted for the various
phases of the light-curve (‘plateau’, ‘shallow’ and ‘steep’
– defined in Table 2); the results of the fitting are pre-
sented in Table 3. In each phase, the spectrum could
be well modelled by a single power-law (no break re-
quired), with excess absorption in the rest-frame of the
GRB (modelled using ztbabs and the ‘Wilms’ abun-
dance in xspec; Wilms et al. 2000). Together with the
WT spectrum from ∼ 200–590 s after the trigger (in the
plateau stage), the first two PC spectra (plateau and
shallow) are fully consistent with a constant photon in-
dex of Γ = 2.07 ± 0.06 and NH,z = (9.2 ± 1.2) × 10
cm−2.
Following the second apparent break in the light-curve,
around 3.2 × 104 s, the spectrum hardened slightly, to a
photon index of Γ = 1.83 ± 0.11 (or 1.87 ± 0.08 using
NH,z = 9.2 × 10
21 cm−2).
2.2.2. XMM-Newton
XMM-Newton (Jansen et al. 2001) performed a Target
of Opportunity observation of GRB 061121 (Observation
ID 0311792101) less than 6.5 hr after the trigger (Schartel
2006) and collected data for ∼ 38 ks (MOS1, MOS2;
Turner et al. 2001) and ∼ 35 ks (PN; Strüder et al. 2001).
This observation is mainly during the ‘shallow’ phase,
though also covers a short timespan after the break at
around 32 ks.
Figure 6 plots the PN flux light-curve and hardness
ratio during the XMM-Newton observation, showing the
lack of spectral evolution during this time frame; a hard-
ness ratio calculated for the Swift data was in agreement
with this finding. The decay slope over this time (MOS1,
MOS2, PN and joint) is consistent with the Swift results
(α ∼ 1.3; note this crosses the time of the second break
in the decay).
The XMM-Newton EPIC29 spectra show clear evidence
for excess NH, in agreement with the Swift data. In addi-
tion, fitting with excess NH in the rest-frame of the GRB
29 European Photon Imaging Camera
8 K.L. Page et al.
Epoch time since Γ NH,z χ
2/ν corresponding
trigger (s) (1021 cm−2) α
Plateau 590–1560 2.14 ± 0.12 10.8
62.5/52 0.38 ± 0.08
Shallow 4900–22245 2.04 ± 0.10 8.9
67.5/70 1.07
+0.04
−0.06
Steep 34550–1152750 1.83 ± 0.11 8.0
48.0/55 1.53
+0.09
−0.04
Plateau 590–1560 2.09 ± 0.08 9.2 ± 1.2 (tied) 63.5/53 0.38 ± 0.08
Shallow 4900–22245 2.05 ± 0.06 9.2 (tied) 67.6/71 1.07
+0.04
−0.06
Steep 34550–1152750 1.87 ± 0.08 9.2 (tied) 48.7/56 1.53
+0.09
−0.04
TABLE 3
XRT PC spectral fits - rest-frame NH free and then tied between all three spectra. The temporal decay slopes, α,
corresponding to each stage are also given. The Galactic absorbing column of NH = 5.09 × 10
20 cm−2 was always
included in the model.
3×104 3.5×104 4×104 4.5×104 5×104 5.5×104 6×104
time since BAT trigger (s)
Fig. 6.— XMM-Newton EPIC-PN light-curve and hardness ratio
of GRB 061121. The horizontal line shows the hardness ratio is
consistent with a constant value of ∼ 1.46, indicating there is no
spectral evolution during this time.
gives a significantly better fit than at z = 0, as shown in
Figure 7. When fitting in the observer’s frame there is a
noticeable bump in the residuals around 0.6 keV; fitting
with NH at z = 1.314 removes this feature. The data
are of sufficiently high signal-to-noise that the redshift of
the absorber can be estimated from the spectrum. Limits
can be placed on the redshift and absorbing column, re-
spectively, of z > 1.2 and NH,z > 4.6 × 10
21 cm−2 at 99%
confidence, in agreement with the spectroscopic redshift
from Bloom et al. (2006) within the statistial uncertain-
ties. At their value of z = 1.314, the excess NH,z from
the EPIC-PN spectrum is (5.3 ± 0.2)× 1021 cm−2, lower
than the best fit to the Swift data from the simultaneous
‘shallow’ decay section, but more similar to the values ob-
tained from fitting the optical-to-X-ray Spectral Energy
Distributions (SEDs) in §2.4.2. In agreement with the
simultaneous XRT PC mode data, there is no evidence
for a break in the EPIC spectrum over this time period.
Spectra from neither the Reflection Grating Spectrom-
eter (den Herder et al. 2001) nor EPIC show obvious
absorption or emission lines.
2.2.3. Chandra
Chandra performed a 33 ks Target of Opportunity ob-
servation at ∼ 61 day after the trigger. No source was
detected at the position of the X-ray afterglow, with a
3σ upper limit of 2.5 × 10−15 erg cm−2 s−1.
2.3. Optical/UV Data
zNH = 5.3x10
21 cm−2
10.5 2 5
channel energy (keV)
NH = 1.3x10
21 cm−2
Fig. 7.— EPIC-PN spectrum of the late-time afterglow of
GRB 061121, with an excess absorbing column both in the rest-
frame of the GRB and the observer’s frame. The spectrum is much
better modelled with an excess column at z = 1.314.
2.3.1. UVOT
The UVOT detected an optical counterpart in the ini-
tial White filter30 observation, starting 62 s after the
trigger, and subsequently in all other filters (optical and
UV). The UVOT followed the typical sequence for GRB
observations, with the early data being collected in event
mode, which has a frame time of 8.3 ms during this ob-
30 The White filter covers a broad bandpass of λ ∼ 1600−6500 Å.
GRB 061121: Broadband observations 9
servation.31 Photometric measurements were obtained
from the UVOT data using a circular source extraction
region with a 5− 6′′ radius. uvotmaghist was used to
convert count rates to magnitudes and flux; no normali-
sation between the different filters was applied.
As in the gamma-ray and X-ray bands, the main burst
was detected, with an increase in count rate seen between
∼ 50–75 s after the trigger (Figures 1 and 4). However,
although an increase in count rate is seen for the UVOT
data, it is by a smaller factor than observed for the XRT.
After ∼ 110 s, the UVOT emission stops decaying and re-
brightens slightly, until 140 s after the trigger, at which
time it flattens off and then starts to fade again (Fig-
ure 4). The slower decay between ∼ 100–200 s may be
indicative of the contribution of an additional (afterglow)
component beginning to dominate.
A single UV/optical light curve was created from all
the UVOT filters in order to get the best measurement of
the optical temporal decay. This was done by fitting each
filter dataset individually (between 200 and 1 × 105 s)
and finding the normalisation, which was then modified
to correspond to that of the V -band light-curve. The
decay across all the filters beyond 200 s after the trigger
can be fitted with a single slope of αUVOT = 0.68 ± 0.02;
the individual U , B and V decay rates are consistent
with one another. No break in the light-curve is seen out
to ∼ 100 ks.
2.3.2. ROTSE
ROTSE-IIIa, at the Siding Spring Observatory in Aus-
tralia, first imaged GRB 061121 21.6 s after the trigger
time under poor (windy) seeing conditions. A variable
source was immediately identified, at a position coinci-
dent with that determined by the UVOT (Yost et al.
2006).
The ROTSE data (unfiltered, but calibrated to the R-
band) have been included in Figure 10 (discussed later).
It is noticeable that the peak around 75 s seen in the Swift
data is not readily apparent in the ROTSE measure-
ments. The bandpass of the UVOT White filter is more
sensitive to photons with wavelengths of λ < 4500Å32,
while the ROTSE bandpass is redder. This, together
with poor seeing conditions during the observation, may
explain why the ROTSE light-curve does not clearly show
the main emission.
2.3.3. Faulkes Telescope North
The FTN, at Haleakala on Maui, Hawaii, began ob-
servations of GRB 061121 225 s after the burst trigger,
performing a BV Ri′ multi-colour sequence (Melandri
et al. 2006). R-band photometry was performed rela-
tive to the USNO-B 1.0 ‘R2’ magnitudes. Magnitudes
were then corrected for Galactic extinction using the
dust-extinction maps by Schlegel et al. (1998), and con-
verted to fluxes using the absolute flux calibration from
Fukugita et al. (1995). The photometric R-band points
have been included in Figure 10.
31 The data have been adjusted to take into account an incor-
rect onboard setting (between 2006-11-10 and 2006-11-22), which
resulted in the wrong frame times being stored in the headers of
the UVOT files (Marshall 2006a).
32 See http://swiftsc.gsfc.nasa.gov/docs/heasarc/caldb/swift/
2.4. Broadband Modelling
2.4.1. Gamma-rays – X-rays
Spectral Analysis— Because the BAT was in event
mode throughout the observation of the main burst of
GRB 061121, detailed spectroscopy could be performed.
Unfortunately this was not the case during the prompt
observation of GRB 060124 (Romano et al. 2006).
Figure 2 demonstrates the spectral evolution seen in
both the BAT and XRT during the prompt emission.
Spectra were extracted over 2 s intervals, in an attempt
to obtain sufficient signal to noise while not binning over
too much of the rapid variability. The BAT data are
hardest around 68 s and 75 s (the second of these times
corresponding to the peak of the main emission); the
XRT hardness peaks about 70 s, which could be a further
indication of the softer data lagging the harder. The joint
spectrum (Γjoint comes from a simple absorbed power-
law fit to the simultaneous BAT and XRT data) is at its
hardest during the brightest part of the emission. The
joint fit also hardens around 68–70 s, between the times
when the BAT and XRT data respectively are at their
hardest. The onboard spectral time-bin selection pre-
vents the Konus-Wind data from being sliced into corre-
sponding times, so constraints have not been placed on
the high energy cut-off, Epeak. Breaks in the XRT-BAT
power-laws can only be poorly constrained.
In Figure 4, the BAT and XRT data were converted to
0.3–10 keV fluxes using the time-sliced power-law fits to
the simultaneous BAT and XRT spectra. Without the
use of such varying conversion factors, the derived BAT
and XRT fluxes would have been inconsistent with each
other.
A broadband spectrum, covering 0.3 keV to 4 MeV in
the observer’s frame (XRT, BAT and Konus-Wind) for
∼ 70–75 s post trigger was fitted by the absorbed cut-off
power-law model described in §2.1.2. A constant factor
of up to 10% was included between the BAT and Konus-
Wind data, to allow for calibration uncertainties. The
best fit (χ2/dof = 301/167) gives Γ = 1.19 ± 0.01, with
Epeak = 670
−47 keV. NH,z was fixed at 9.2 × 10
21 cm−2
(from the X-ray fits in §2.2.1). Allowing Γ to vary be-
tween the three spectra hints at further spectral curva-
ture, although the differences are marginal, significant at
only the 2σ level.
The isotropic equivalent energy (calculated using
the time-integrated flux over the full T90 period) is
2.8 × 1053 erg in the 1 keV – 10 MeV band (GRB rest
frame), meaning that GRB 061121 is consistent with the
Amati relationship (Amati et al. 2002). See §3.3.2 for a
beaming-corrected gamma-ray energy limit.
Lag Analysis— A lag analysis (e.g., Norris et al. 1996)
between the BAT bands leads to interesting results.
Comparing bands 50–100 keV and 15–25 keV, the precur-
sor emission yields a spectral lag of 600 ± 100 ms, while
the main emission has a much smaller lag of 1 ± 6 ms.
Note that the calculation was performed using 64 ms
binning for the precursor and 4 ms binning for the main
burst; see Norris (2002) and Norris & Bonnell (2006) for
more details on the procedure. This lag for the main
emission is rather small for a typical long burst, however
both lags are consistent with the long-burst luminosity-
lag relationship generally seen (Norris et al. 2000). The
http://swiftsc.gsfc.nasa.gov/docs/heasarc/caldb/swift/
10 K.L. Page et al.
0 2 4 6
time delay (s)
1−4 keV
4−10 keV
15−25 keV
25−50 keV
50−100 keV
100−150 keV
Fig. 8.— Autocorrelation function of the BAT and XRT data
during the prompt emission of GRB 061121, showing that the main
burst peak is broader at softer energies.
short spectral lag for the main emission, and the longer
value for the precursor are also found when comparing
the 100–350 keV and 25–50 keV bands.
Similarly, comparison of the hard and soft (2–10 keV
and 0.3–2 keV) XRT bands reveals a lag of approximately
2.5 s, as the emission softens through the main burst.
The X-ray data also lag behind the gamma-ray data,
and the optical behind the X-ray.
Link et al. (1993) and Fenimore et al. (1995) used a
sample of BATSE33 (Paciesas et al. 1999) bursts to in-
vestigate the relationship between the duration of bursts
and the energy band considered. They found that the
bursts, and smaller structures within the main emission,
generally become shorter with increasing energy (see also
Cheng et al. 1995; Norris et al. 1996; in’t Zand & Fen-
imore 1996; Piro et al. 1998). Figure 8 plots the auto-
correlation function over various X-ray and gamma-ray
bands, to reinforce the point that the peak is narrower
the harder the band – over X-ray as well as gamma-ray
energies. Comparison of the light-curves over the differ-
ent energy bands in Figure 1 demonstrates this as well. A
similar behaviour was also found for GRB 060124, where
Romano et al. (2006) compared the T90 values obtained
for the main burst over the X-ray and gamma-ray bands.
Fenimore et al. (1995) found that the width of the auto-
correlation function, W ∝ E−0.4, where E is the energy
at which the function was determined; the six measure-
ments from GRB 061121 are consistent with this finding.
2.4.2. Optical – X-rays
Using the Swift X-ray and UV/optical data, R and
i′ band data from the Faulkes Telescope and Rc data
from the Kanata telescope (Uemera et al. 2006), SEDs
were produced at epochs corresponding to the peak of
the emission (72–75 s post BAT trigger), the plateau
stage and during the shallow decay. Fitting at the differ-
ent epochs gives an estimation of the broadband spectral
variation.
For each of the UVOT lenticular filters, the tool
uvot2pha was used to produce spectral files compatible
with xspec, and for the latter two epochs the count rate
33 Burst And Transient Source Experiment
in each band was set to that determined from a power-law
fit to the individual filter light curves over the time inter-
val in question, using α = 0.68. To determine the Faulkes
Telescope R and i′ band flux during the plateau stage, a
power law was fitted to the complete data set (220–1229 s
post BAT trigger for R and 467–1401 s for i′) with the
decay index left as a free parameter. The R magnitude
at the mid-time of the shallow stage (6058 s) was deter-
mined from the Kanata R-band magnitude reported at
6797 s (Uemera et al. 2006), assuming the same decay
index as observed in the UVOT data. An uncertainty of
0.2 mag was assumed as the systematic uncertainty for
the photometric calibration of the ground based data.
At a redshift of z = 1.314, the beginning of the Lyman-
α forest is redshifted to an observer-frame wavelength of
∼ 2812 Å which falls within the UVW1 filter bandpass,
the reddest of the UV filters. A correction was applied to
the three UV filter fluxes to account for this absorption,
based on parameters from Madau (1995) and Madau et
al. (1996); see also Curran et al. (in prep).
The methods used for simultaneous fitting of the SED
components are described in detail in Schady et al.
(2007a). The SEDs were fitted with a power-law, or
a broken power-law, as expected from the synchrotron
emission, and two dust and gas components, to model
the Galactic and host galaxy photoelectric absorption
and dust extinction. The column density and reddening
in the first absorption system were fixed at the Galactic
values. [The Galactic extinction along this line of sight
is E(B − V ) = 0.046 (Schlegel et al. 1998).]
The second photoelectric absorption system was set to
the redshift of the GRB, and the neutral hydrogen col-
umn density in the host galaxy was determined assuming
Solar abundances. The dependence of dust extinction on
wavelength in the GRB host galaxy was modelled us-
ing three extinction laws, taken from observations of the
Milky Way (MW), the Large Magellanic Cloud (LMC)
and the Small Magellanic Cloud (SMC) and parame-
terised by Pei (1992) and Cardelli et al. (1989). The
greatest differences observed in these extinction laws are
the amount of far UV extinction (which is greatest in
the SMC and least in the MW) and the strength of the
2175 Å absorption feature (which is most prominent in
the MW and negligible in the SMC).
Fitting these data together, a measurement of the spec-
tral slope and optical and X-ray intrinsic extinctions (for
the second two epochs) were obtained (Table 4); the AV
values given in the table are in addition to the AV = 0.151
associated with the Milky Way itself. The slope above
the break energy (which lies towards the low energy end
of the X-ray bandpass for each phase) was assumed to
be exactly 0.5 steeper than the spectral slope below the
break (the condition required for a cooling break), since
allowing all of the parameters to vary leads to uncon-
strained fits. Figure 9 shows, as an example, the fit to
the data in the plateau stage.
A Milky Way dust extinction law provides the best
overall fit to the data, using a broken power-law model,
although the LMC model is equally acceptable.
During the plateau phase, and adopting the bro-
ken power-law model parameters given in Table 4, we
find gas-to-dust ratios of (1.6 ± 0.7), (2.6 ± 0.7) and
(3.0 ± 0.7) ×1022 cm−2 mag−1 for MW, LMC and SMC
GRB 061121: Broadband observations 11
1015 1016 1017 1018
Frequency (Hz)
Fig. 9.— Broken power-law fit to the UVOT, XRT and ground-
based R and i′ spectral energy distribution of GRB 061121 between
∼ 596–1566 s after the trigger (plateau phase) plotted in the ob-
server’s frame. The arrows indicate the beginning of the Lyman-α
forest (1215Å in the rest-frame) and the absorption feature in the
MW dust extinction law (2175Å), which is shown by a dotted line.
The solid line corresponds to the LMC extinction, and the dashed
one to the SMC extinction.
fits respectively. We can compare these estimates to the
measured values for the MW of (4.93 ± 0.45) × 1021
cm−2 mag−1 (Diplas & Savage 1994) and the LMC and
SMC of (2.0 ± 0.8) and (4.4 ± 1.1) ×1022 cm−2 mag−1 ,
respectively (Koornneef 1982; Bouchet et al. 1985). The
MW fit to GRB 061121, which is found to be marginally
the best model, is consistent with the LMC gas-to-dust
ratio only, at the 90% confidence level. The ratios derived
from the LMC and SMC fits are consistent with both the
LMC and SMC gas-to-dust ratios. We note that all fits
are inconsistent with the MW ratio at this confidence
level, following the trend seen in pre-Swift bursts (e.g.,
Starling et al. 2007 and references therein), and that if
a metallicity below Solar were adopted, the gas-to-dust
ratio of GRB 061121 would increase, moving it further
towards the SMC value.
3. DISCUSSION
Swift triggered on a precursor to GRB 061121 leading
to comprehensive broadband observations of the prompt
emission, as well as the later afterglow. We discuss these
here, together with possible mechanisms involved.
3.1. Precursor
Lazzati (2005) found that about 20% of BATSE bursts
showed evidence for gamma-ray emission above the back-
ground between 10 to ∼200 s before the main burst, typi-
cally with non-thermal spectra which tended to be softer
than the main burst. GRB 060124 (Romano et al. 2006)
and GRB 061121 show the same behaviour.
Precursor models have been proposed for emission well-
separated from, or just prior to, the main burst. Early
emission occurring only a few seconds before the main
burst has been explained by the fireball interacting with
the massive progenitor star – though the spectrum of
such emission is expected to be thermal (Ramirez-Ruiz
et al. 2002a). Lazzati et al. (2007) investigated shocks in
a cocoon around the main burst; their model predicts a
non-thermal precursor as the jet breaks out of the sur-
face of the star. A high-pressure cocoon is formed as the
sub-relativistic jet head forces its way out of the star. As
the head of the jet breaks through the surface, the energy
of this cocoon is released through a nozzle and can give
rise to a precursor (Ramirez-Ruiz et al. 2002a,b). Within
the framework of this model, observers located at view-
ing angles of 5◦ < θ < 11◦ are expected to see first a
relatively bright precursor, then a dark phase with lit-
tle emission, followed, when the jet enters the unshocked
phase, by a bright GRB; this is very similar to the light-
curve observed for GRB 061121. Waxman & Mészáros
(2003) demonstrate that both a series of thermal X-ray
precursors (becoming progressively shorter and harder)
and nonthermal emission can be produced by an emerg-
ing shocked jet, although the nonthermal component is
expected to be in the MeV range. There could also be
an accompanying inverse Compton component, formed
by the thermal X-rays being upscattered by the jet.
The same type of smooth, wide-pulse, low intensity
emission as seen in some precursors, but occurring af-
ter the main emission is also occasionally seen (e.g.,
Hakkila & Giblin 2004; Nakamura 2000). Hakkila & Gib-
lin (2004) discuss two examples where postcursor emis-
sion is found to have a longer lag than expected from the
lag-luminosity relation, smoother shape and to be softer.
In the case of the GRB 061121 precursor, the spectrum is,
indeed, softer than the main event, and shows a compar-
atively smooth profile. The emission does have a longer
lag than the main emission, but it is still consistent with
the lag-luminosity relation.
There are two expected effects which could lead to such
a difference in lags for separate parts of a single burst: the
much lower luminosity for the precursor (resulting from
a much smaller Lorentz factor; the measured fluence of
the precursor is about a factor of 30 smaller than the
fluence of the main emission) is a natural explanation,
while the precursor being emitted at a greater off-axis
angle could also have an effect. In this second case, ejecta
are considered to emerge at different angles with respect
to the jet axis; not all of the solid angle of the jet will be
‘filled’ uniformly.
Such late postcursor emission is unlikely to be linked to
the jet breakout from the stellar surface, and it may not
be sensible to attribute apparently similar phenomena
(in the form of pre- and postcursors) to entirely different
processes.
Pre/postcursor emission could be due to the decelera-
tion of a faster front shell, resulting in slower shells catch-
ing up and colliding with it (Fenimore & Ramirez-Ruiz
1999; Umeda et al. 2005; note, however, that a faster
shell would be inconsistent with the precursor having a
smaller Lorentz factor as suggested to explain the lag
discrepancy), or late activity of the central engine. The
presence of flares in about 50% of Swift bursts is gener-
ally attributed to continuing activity of the central en-
gine (Burrows et al. 2005b; Zhang et al. 2006a) and the
appearance of broken power-laws in the X-ray spectra of
12 K.L. Page et al.
X-ray Model Extinction NH,z Γ1 Ebreak Γ
χ2/dof
Epoch (1021 cm−2) (keV)
Peak PL SMC 1.6
0.99 ± 0.01 · · · · · · 0.64 25/27
LMC 1.9
1.06 ± 0.01 · · · · · · 0.98 23/27
MW 2.4
1.16 ± 0.01 · · · · · · 1.51 22/27
BKN PL SMC 2.7+9.5
0.72+0.08
−0.15
0.17+0.79
−0.15
1.22 0.51 22/26
LMC 3.0+7.7
0.77+0.08
−0.20
0.18+0.53
−0.17
1.27 0.72 22/26
MW 3.0+7.7
0.77+0.10
−0.21
0.09+0.30
−0.09
1.27 1.03 22/26
Plateau PL SMC 1.42± 0.51 1.58± 0.02 · · · · · · 0.62± 0.05 167/59
LMC 1.98± 0.54 1.64± 0.03 · · · · · · 0.94± 0.08 152/59
MW 2.71± 0.69 1.71± 0.03 · · · · · · 1.39± 0.10 136/59
BKN PL SMC 3.89
+0.72
−1.01
+0.03
−0.02
+0.36
−0.12
1.96 0.52± 0.04 84/58
LMC 4.40
+0.77
−1.30
+0.04
−0.02
+0.16
−0.14
2.01 0.74± 0.06 80/58
MW 3.91
+0.77
−0.75
+0.04
−0.03
+0.25
−0.20
2.08 1.03
+0.09
−0.08
79/58
Shallow PL SMC 2.72± 0.49 1.69± 0.02 · · · · · · 0.65± 0.04 162/77
LMC 3.37
+0.53
−0.49
1.75± 0.03 · · · · · · 0.98
+0.07
−0.06
146/77
MW 4.60
+0.65
−0.60
1.87± 0.04 · · · · · · 1.63
+0.12
−0.11
127/77
BKN PL SMC 4.02+0.62
−0.67
1.58+0.02
−0.03
1.30+0.19
−0.11
2.08 0.50± 0.04 101/76
LMC 4.41+0.69
−0.63
1.62± 0.03 1.30+0.16
−0.14
2.12 0.72± 0.06 99/76
MW 4.78+0.75
−0.65
1.67± 0.04 1.35+0.16
−0.17
2.17 1.02+0.11
−0.10
102/76
a Γ2 is set to be equal to Γ1 + 0.5 in each broken power-law fit, as would be expected if the change in index were due to a cooling
break.
b In the fit to the peak epoch, AV is fixed to the average best-fit value found in the same model fits to plateau and shallow stage
data. The AV values are given for the observer’s frame of reference.
TABLE 4
Power-law (PL) and broken power-law (BKN PL) fits to the simultaneous UVOT and XRT spectra of GRB 061121, for
three different dust extinction models: Small and Large Magellanic Clouds (SMC and LMC) and the Milky Way (MW).
Γ1 and Γ2 are the photon indices below and above the spectral break for the BKN PL models. The data points have not
been corrected for reddening.
both flares and the prompt emission (Guetta et al. 2006;
Goad et al. 2007) hints of a common mechanism.
3.2. Prompt Emission
The prompt emission mechanism for GRBs is still de-
bated and the origin of Epeak is not fully understood
(Mészáros et al. 1994; Pilla & Loeb 1998; Lloyd & Pet-
rosian 2000; Zhang & Meszaros 2002; Rees & Mészáros
2005; Pe’er et al. 2005). The standard synchrotron
model predicts fast cooling (Ghisellini et al. 2000) with
a photon index, Γ, of 3/2 and (p/2)+1 below and above
the peak energy, respectively (e.g., Zhang & Mészáros
2004). The Konus-Wind spectral index below Epeak is
shallower than 3/2, which may suggest a slow cooling
spectrum with p < 2 [Epeak being the cooling frequency
and Γ =(p+1)/2] or additional heating. A slow-cooling
spectrum can be retained by assuming that the magnetic
fields behind the shock decay significantly in 104–105 cm,
so that synchrotron emission happens in small scale mag-
netic fields (Pe’er & Zhang 2006).
The SED at the peak time (SED 2 in Figure 11,
discussed below) has a peak flux density of around
1 keV, below which the optical to X-ray spectral slope
is 0.11 ± 0.09. This slope is harder than expected from
the standard synchrotron model (which predicts an in-
dex of 1/3). There should, however, be spectral cur-
vature around the break, which could flatten the index
(Lloyd & Petrosian 2000), so the data could still be con-
sistent with the synchrotron model. An alternative to
synchrotron emission, in the form of ‘jitter’ radiation is
discussed by Medvedev (2000), though that model pre-
dicts an even steeper index of 1 below the jitter break
frequency.
Figures 4 and 10 show that all three instruments on-
board Swift saw the prompt emission around 75 s after
the BAT trigger. However, it is noticeable that most of
the emission is in the gamma-ray and X-ray bands, with
the optical showing a relatively small increase in bright-
ness in comparison. Assuming the observed process is
synchrotron, then the prompt emission which is detected
by the UVOT will be the low-frequency extension of this
in the internal shock. No reverse shock is apparent.
3.3. Afterglow Emission
3.3.1. Broken Power-law Decline Models
The afterglow of GRB 061121 was observed over an
even broader energy range (from radio to X-rays) than
the prompt emission, with multi-colour data being ob-
tained from ∼ 100–105 s after the trigger. The X-ray
light-curve shows evidence for substantial curvature at
later times (see Figure 3), as has been found for other
Swift GRBs (e.g., GRBs 050315 – Vaughan et al. 2006;
060614 – Gehrels et al. 2006b). The standard practice
has been to fit such a decay using a series of power-
law segments as a function of time. An alternative
exponential-to-power-law description of the light-curve
is given in §3.3.2.
Nousek et al. (2006) and Zhang et al. (2006a) have
both discussed the canonical shape that many Swift af-
terglows seem to follow: steep to plateau to shallow,
with some light-curves showing a further steepening. In
these previous works, the extrapolation of the BAT data
into the XRT band was incorporated into the derivation
of the steep decay at the start of the canonical light-
curve shape. In the case of GRB 061121, the full curve
GRB 061121: Broadband observations 13
can be seen entirely in X-rays, suggesting that the pre-
vious extrapolations are reliable. For the afterglow of
GRB 061121, only data after the end of the main burst
have been modelled with power-laws. The early steep
decline, which might be attributable to the curvature ef-
fect (Kumar & Panaitescu 2000; Dermer 2004; Fan &
Wei 2005), is not considered here.
According to the model proposed in Nousek et al.
(2006) and Zhang et al. (2006a), the plateau phase of
the light-curve is due to energy injection in the fireball.
The plateau phase of GRB 061121 is consistent with an
injection of energy since the luminosity index, q, is nega-
tive, which is the requirement for injection to modify the
afterglow (Zhang et al. 2006a); the later two stages both
have q > 1. However, as will be discussed in §3.3.2, the
plateau and final transition to the power-law decay are
only visible in the X-ray data for GRB 061121; the start
of the final decay is much earlier in the V and R-bands
(see Figure 10). One might expect that energy injection
would affect all the energy bands simultaneously, rather
than just the X-rays.
From the standard afterglowmodel computations (e.g.,
Zhang & Mészáros 2004), we find that none of the closure
relations fit the entire dataset completely: although the
shallow phase (after the end of energy injection, between
T + 2.3 ks and T + 32 ks) could be consistent with the
evolution of a blast-wave which had already entered the
slow cooling regime when deceleration started [i.e., ν >
max(νm, νc) where νc is the cooling frequency and νm
is the synchrotron injection frequency; Sari et al. 1998;
Chevalier & Li 2000], the steeper part of the decay curve
(T> 32 ks) is not consistent with any of the models. This
lack of consistency suggests that a different approach is
required.
The change in decay slope between the shallow to
steep phases (∼ 32 ks) cannot be easily identified with
a jet-break. It certainly seems unlikely that the simplest
side-spreading jet model could be applicable, since the
post-break decay index (α ∼ 1.5) is not steep enough
(a post-jet decay has α = p, where p is the electron in-
dex). There is also some indication that the X-ray spec-
tral slope hardens after the break, whereas no change
in spectral signature is expected over a jet-break. In
the case of a non-laterally expanding jet (Panaitescu
& Mészáros 1999), α = (3β/2) + 0.25 [for a homoge-
neous circumstellar medium (CSM); Panaitescu et al.
(2006)], which does, indeed, fit the data after this break:
[1.5 × (0.9 ± 0.08)] + 0.25 = 1.6 ± 0.1; the measured
α is 1.53. Such a confined jet has been suggested as an
explanation for the observed decay in a number of pre-
vious bursts (e.g., GRB 990123 – Kulkarni et al. 1999;
GRB 050525A – Blustin et al. 2006; GRB 061007 –
Schady et al. 2007b). The UVOT data obtained around
this time show little evidence for a break, whereas jet
breaks should occur across all energy bands simultane-
ously. However non-simultaneity could be explained by
a multi-component outflow, where the X-ray emission is
produced within a narrow jet, while the optical compo-
nent comes from a wider jet with lower Lorentz factor
(Panaitescu & Kumar 2004; Oates et al. 2007). There
remains the issue, however, that α should steepen by
0.75 over a jet break (Mészáros & Rees 1999), whereas
the maximum observed change (within the 90% errors)
is only ∆α < 0.61, excluding ∆α = 0.75 at almost 3σ;
also, again there should be no spectral evolution across
the break. There is, however, a probable jet break at
later times, which will be covered in the next Section.
Other multi-component models [see, e.g., Oates et al.
(2007) and references therein] also fail to explain the
data, because of the lack of observed energy injection
(plateau phase) in the optical data.
Panaitescu et al. (2006a) discuss chromatic breaks in
Swift light-curves, and postulate that these could be due
to a change in microphysical parameters within a wind
environment. However, this model requires the cooling
frequency to lie between the X-ray and optical bands
and, as will be discussed in §3.3.2, this does not seem to
be the case here.
3.3.2. Exponential-to-power-law Decline Model
As first described by O’Brien et al. (2006), and further
expanded by Willingale et al. (2007), GRB light-curves
can be well modelled by one or two components com-
prised of an early exponential rise followed by a power-
law decay phase. Of these components, the first repre-
sents the prompt gamma-ray emission and early X-ray
decay. The second, when detected, dominates at later
times, forming what we see as the afterglow. These re-
sults show that fitting an intrinsically curved decay with
multiple power-law segments runs the risk of incorrectly
identifying temporal breaks (see also Sakamoto et al. in
prep). In this Section the models of O’Brien et al. (2006)
and Willingale et al. (2007) are applied to the multi-band
afterglow data of GRB 061121.
Figure 10 brings together the BAT, XRT, UVOT, FTN
and ROTSE data, along with further optical and ra-
dio points taken from the GCN Circulars (Halpern et
al. 2006a,b; Halpern & Armstrong 2006a,b; Chandra &
Frail 2006; van der Horst et al. 2006a,b) and the upper
limit from Chandra, to form a multi-energy decay plot.
The data have been plotted as ‘time since trigger + 4 s’
in order to include the precursor on a log time-scale. The
optical points have all been corrected for extinction using
AV = 1.2 (a combination of the Galactic value of 0.151
and an estimate of AV ∼ 1 for the GRB host galaxy –
see §2.4.2).
The contribution from the host galaxy reported by
Malesani et al. (2006) and Cobb (2006) has been sub-
tracted from the V - and R-band flux values. The magni-
tude of the host in the V -band is 22.4, which only changes
the last two or three V -band points by a small amount.
For the R-band we have no direct measurement, but the
last group of MDM exposures gave an R magnitude of
22.7, corresponding to a flux level of 2.8 µJy, and the flux
level is still declining at that epoch (∼ 3.3 × 105 s), so
an R-band flux level of 2.5 µJy was adopted for the host.
The error bars shown on the last few points reflect the
large uncertainty in the galaxy contribution subtracted.
The curved dotted lines in Figure 10 are the fits to
the data using the exponential-to-power-law model, fol-
lowed by a break to a steeper decay around 105 s. These
models are parameterised by the power-law decay, α,
and Ta, the time at which this decay is established.
For the X-ray data, Ta,X is found to be 5250
−460 s
and αa,X = 1.32 ± 0.03. Fits were also performed to
the V - and R-band data, yielding: αa,V = 0.66 ± 0.04
(with Ta,V = 70
−70 s) and αa,R = 0.84 ± 0.03
14 K.L. Page et al.
Fig. 10.— Flux density light-curves for the gamma-ray, X-ray, optical and radio data obtained for GRB 061121. The vertical dotted lines
indicate the times used for the SED plots shown in Figure 11, while the curved dotted lines show the fit to the X-ray and optical data,
including a late-time break, as described in the text.
Fig. 11.— SEDs for the four time intervals indicated in Figure 10.
SED 2 (the peak of the burst emission) includes the Konus-Wind
data, although these have not been included in Figure 10. The
solid lines represent the power-law fits to the BAT, XRT and Konus
data, while the dashed lines join the radio, optical and 1 keV points.
Spectral evolution over time is clearly seen.
(Ta,R = 230
−230 s).
The non-detection by Chandra almost two months af-
ter the burst shows there must have been a further steep-
ening in the X-ray regime, and the optical data are not
inconsistent with this finding. Constraining the tempo-
ral index after the late break to be α = 2 (a typical slope
for a post-jet-break decay), break times of ∼ 2.5 × 105,
∼ 2.5 × 104 and ∼ 105 s are estimated for the X-ray, V -
and R-band respectively; note that the UVOT V -band
value is particularly uncertain, given the small number
of data points at late times. Within the uncertainties,
these times are likely to be consistent, so the turnover
could be achromatic, as required for a jet break. From
Willingale et al. (2007), a jet break might be expected at
∼ 100 × Ta,X – i.e., 5.5 × 10
5 s, which is in agreement
with these fits.
As can be seen from these numbers and the mod-
els plotted in Figure 10, the X-ray data clearly show
the transition from the plateau to the power-law decay,
whereas the start of the final decay is much earlier in the
V - and R-bands. The V -band decay is also significantly
flatter (by α ∼ 0.2) than that estimated for the R-band.
As previously stated, the V , B and U light-curves are
all consistent with this slow decay. There have been few
multi-colour optical decay curves obtained for GRB af-
terglows, and, of these, the different filters [in the case of
GRB 061007 (Schady et al. 2007; Mundell et al. 2007)
X-ray and gamma-ray data as well as the optical] tend
to track each other (e.g., Guidorzi et al. 2005; Blustin
et al. 2006; de Ugarte Postigo et al. 2007). In the case
of GRB 061121, we find that the R-band data are fading
more rapidly than the V . GRB 060218, which was as-
sociated with a supernova (e.g., Campana et al. 2006a),
shows changes throughout the optical spectra, because of
a combination of shock break-out and radioactive heat-
ing of the supernova ejecta. There is a large difference
between the decays of the blue (V , U , B) and red (R)
data for GRB 061121, which cannot be easily explained
by a synchrotron spectrum. Although no supernova has
been detected in this case, we speculate that some form
of pre-supernova thermal emission could possibly be af-
fecting the optical data, adding energy into the blue end
of the spectrum, thus slowing its decline.
After the break in the decays around 105 s, the light-
curves across all bands become more consistent with one
GRB 061121: Broadband observations 15
another, although there are only limited data at such a
late time.
The vertical dotted lines in Figure 10 show the times
of the SEDs plotted in Figure 11; again, all points were
corrected for an extinction of AV = 1.2, so that they
represent the true SEDs (with the frequency in the ob-
server’s frame). The solid lines represent actual fits to
the X-ray and gamma-ray data, while the dashed lines
just join the separate radio, optical and 1 keV points.
The times of these SEDs, which clearly show spectral
evolution, correspond to (1) before the main BAT peak,
56 s after trigger; (2) at the BAT peak, 76 s after trig-
ger; (3) just after the start of the plateau, 300 s after the
trigger; (4) in the main decay at 65 ks (chosen because
radio measurements were taken at this time). SEDs 3
and 4 do not contain any BAT or Konus data, since the
gamma-ray flux had decayed by this point; the highest
energy point in these corresponds to the maximum en-
ergy (10 keV) of the X-ray fits.
Table 4 demonstrates that the optical and X-ray spec-
tra during the peak emission are best fitted with a broken
power-law model, with the break energy at the very low
energy end of the X-ray bandpass. SED 2 in Figure 11
shows that this spectral break corresponds to the peak
frequency in a flux density plot (β1 is less than zero in
this case). Only during SED 2 is the optical flux density
lower than that of the higher energy data. Figure 4 also
shows that the optical emission is less strong than the
X-ray and gamma-ray data during the main burst.
Table 5 shows the values of α for the X-ray and opti-
cal decays (i.e., before and after the break) in SED 4, at
65 ks, with their corresponding spectral indices. For the
initial stages of the power-law decay (Ta < t < 65000 s)
the evolution of the afterglow SED and the coupling be-
tween the temporal and spectral indices are not com-
pletely consistent with the standard model: although the
R-band decay, with αa,R = 0.84 ± 0.03, is in good agree-
ment with the homogeneous CSM model below the cool-
ing break, the X-ray and V -band flux decays are slower
than expected from the measured spectral indices; they
are in best agreement with the same constant density
model below νc, however.
The point at which the power-law decay dominates
the exponential in the optical bands is noticeably earlier
than in the X-ray (< few hundred seconds, rather than
∼ 5000 s) and, as mentioned above, the decay indices
are significantly different for all three (X-ray, V and R)
bands (see Figure 10). At the time of SED 3, the X-ray
data are not decaying (i.e., this is during the plateau),
yet both the V and R-band data have already entered
the power-law decline phase. The R-band is decaying
faster than the V -band, so the spectral index through
the optical range is becoming harder. The X-ray spec-
tral index shows a similar hardening trend (see Table 3),
so the SED measured from optical to 10 keV is gradually
getting harder. Such spectral hardening from the plateau
to the final decay is a feature of many X-ray afterglows
(Willingale et al. 2007).
This slow hardening of the broadband spectrum with
time could be a signature of synchrotron self-Compton
emission (Sari & Esin 2001; Panaitescu & Kumar 2000).
The strength of the self-Compton component in the af-
terglow depends on the flux of low energy photons (radio-
optical) and the electron density in the shock. Using the
formulation in Sari & Esin (2001) the density required is
given by
n1 = 3× 10
f ICmax
(E52tday)
−1/3cm−3 (1)
where f ICmax/f
max is the ratio of the peak flux of the
seed synchrotron spectrum (i.e., the source of low energy
photons) and the peak flux of the self-Compton emis-
sion; E52 is the isotropic burst energy in units of 10
52 erg;
tday is the time in days after the burst (which determines
the distance through the CSM swept up by the external
shock). From Figure 11 (SEDs 1, 3 and 4) we see that
f ICmax/f
max ∼ 0.001 if the X-ray flux has a significant con-
tribution from a self-Compton component at tday = 0.75.
A value of E52 = 30 gives n1 ≈ 10
5 cm−3. Even as-
suming the emission at 0.75 days is not dominated by
the self-Comptonisation, and so taking the f ICmax/f
ratio to be a factor of ten smaller, the density would be
∼ 5× 103 cm−3, which is still high. It seems unlikely that
self-Compton emission is the cause of the spectral hard-
ening of the SED unless the CSM density encountered
by the external shock is extremely large. However, there
have been suggestions that GRBs may form in molecular
clouds (Galama & Wijers 2001; Campana et al. 2006b,c),
which have densities of 104 or more particles per cubic
centimetre in the cores (Miyazaki & Tsuboi 1999; Wil-
son et al. 1999). Typically one might expect greater red-
dening than is found here (Table 4), though Waxman &
Draine (2000) discuss the possibility of dust destruction.
The spectrum will be redshifted as the jet slows down,
so the optical and X-ray spectral indices should, if any-
thing, become softer – the opposite of what is seen here.
Although spectral hardening with time is suggested from
the data, it is not be easily explained by current models.
Whether or not there is a Comptonised component,
the later SEDs clearly indicate that there is a break in
the spectrum somewhere between the optical and the X-
ray; this is also shown by the fits in Table 4, where the
UVOT–XRT spectra are better fitted with broken power-
laws, with Ebreak towards the low energy end of the X-ray
bandpass. Since both the optical and X-ray bands ap-
pear to be below the cooling frequency, from the closure
relations given in Table 5, this change in slope cannot
be identified with a cooling break; its origin remains un-
clear.
The redshift of z = 1.314 and the isotropic energy of
Eiso ∼ 3 × 10
53 erg (§2.4.1) can be used to place con-
straints on the jet opening angle. From Sari et al. (1999),
and assuming that the jet break occurs at T0+2× 10
we have θj ∼ 4
)1/8 ( n
where n and ηγ are
the density of the CSM and the efficiency of the fireball
in converting the energy in the ejecta into gamma-rays.
Taking ηγ = 0.2 and n = 3 cm
−3 (following Ghirlanda et
al. 2004), this gives Eγ ∼ 1.7 × 10
51 erg for the beaming-
corrected gamma-ray energy released, which is within the
range previously determined (e.g., Frail et al. 2001) and
consistent with the Ghirlanda relationship (Ghirlanda et
al. 2004).
4. SUMMARY AND CONCLUSIONS
Swift triggered on a precursor to GRB 061121, leading
16 K.L. Page et al.
GRB models α(β) α(βa,X)
α(βopt)
a αbopt
V -band R-band
CSM SCc (νm < ν < νc)
1.49 ± 0.10 1.32 ± 0.03 0.80 ± 0.09 0.66 ± 0.04 0.84 ± 0.03
Wind SCc (νm < ν < νc)
1.99 ± 0.10 1.30 ± 0.09
CSM or Wind SCc & FCd
0.99 ± 0.10 0.30 ± 0.09
(ν > max(νc, νm))
a Decay calculated from the measured spectral index
b Observed power-law decay index.
c Slow cooling.
d Fast cooling.
TABLE 5
Closure relations for exponential-plus-power-law model fits to the X-ray data (βa,X = 0.99± 0.07) and the
optical-to-X-ray band (βopt = 0.53 ± 0.06) from the time of SED 4 (65 ks after the burst).
to unprecedented coverage of the prompt emission by all
three instruments onboard, with the gamma-ray, X-ray
and optical/UV bands all tracking the main peak of the
burst. GRB 061121 is the instantaneously brightest long
Swift burst detected thus far, both in gamma-ray and X-
rays. The precursor and main burst show spectral lags of
different lengths, though both are consistent with the lag-
luminosity relation for long GRBs (Gehrels et al. 2006b).
The SED of the prompt emission, stretching from 1 eV
to 1 MeV shows a peak flux density at around 1 keV and
is harder than the standard model predicts. There is def-
inite curvature in the spectra, with the prompt optical-
to-X-ray spectrum being better fitted by a broken power-
law, similar to results found for fitting X-ray flares (e.g.,
Guetta et al. 2006; Goad et al. 2007).
The afterglow component, in both the optical and
X-ray, starts early on – before, or during, the main
burst peak (see also O’Brien et al. 2006; Willingale
et al. 2007; Zhang et al. 2006b). The broadband
SEDs reveal gradual spectral hardening as the afterglow
evolves, both within the X-ray regime (Γ flattening from
∼ 2.05 to ∼ 1.87) and between the V - and R-band op-
tical data (αV ∼ 0.66 compared with αR ∼ 0.84). Self-
Comptonisation could explain the hardening, although a
molecular-cloud-core density would be required. A prob-
able jet-break occurs around T0 + 2 × 10
5 s, shown by a
late-time non-detection by Chandra. Before this break,
the X-ray and V -band decays are too slow to be readily
explained by the standard models.
This extremely well-sampled burst shows clearly that
there remains much work to be done in the field of GRB
models. A single, unified model for all GRB emission
observed should be the ultimate goal.
5. ACKNOWLEDGMENTS
The authors gratefully acknowledge support for this
work at the University of Leicester by PPARC, at PSU
by NASA and in Italy by funding from ASI. This work is
partly based on observations with the Konus-Wind ex-
periment (supported by the Russian Space Agency con-
tract and RFBR grant 06-02-16070) and on data ob-
tained with XMM-Newton, an ESA science mission with
instruments and contributions directly funded by ESA
Member States and NASA. We thank the Liverpool GRB
group at ARI, Liverpool John Moores University, in par-
ticular C.J. Mottram, D. Carter, R.J. Smith and A.
Gomboc for their assistance with the FTN data acqui-
sition and interpretation. The Faulkes Telescopes are
operated by the Las Cumbres Observatory Global Tele-
scope Network. We also thank J.E. Hill and A.F. Abbey
for discussions and help with the PuPD data, P. Cur-
ran for assistance with the UVOT-XRT SED creation,
D. Grupe for the Chandra upper limit and B. Cobb and
D. Malesani for information regarding the magnitude of
the host galaxy. Thanks as well to C. Akerlof, E. Rykoff,
A. Phillips and M.C.B. Ashley from the ROTSE team.
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