arXiv:0907.3786v2 [astro-ph.GA] 17 Sep 2009
Simultaneous Multi-Wavelength Observations of Sgr A*
during 2007 April 1-11
F. Yusef-Zadeh1, H. Bushouse2, M. Wardle3, C. Heinke4, D. A. Roberts5, C.D. Dowell6, A.
Brunthaler7, M. J. Reid8, C. L. Martin9, D. P. Marrone10, D. Porquet11, N. Grosso11, K.
Dodds-Eden12, G. C. Bower13, H. Wiesemeyer14, A. Miyazaki15, S. Pal16, S. Gillessen12, A.
Goldwurm17, G. Trap18, and H. Maness13
We report the detection of variable emission from Sgr A* in almost all wave-
length bands (i.e. centimeter, millimeter, submillimeter, near-IR and X-rays) during
a multi-wavelength observing campaign. Three new moderate flares are detected
simultaneously in both near-IR and X-ray bands. The ratio of X-ray to near-IR flux
in the flares is consistent with inverse Compton scattering of near-IR photons by
submillimeter emitting relativistic particles which follow scaling relations obtained
from size measurements of Sgr A*. We also find that the flare statistics in near-IR
wavelengths is consistent with the probability of flare emission being inversely pro-
portional to the flux. At millimeter wavelengths, the presence of flare emission at 43
GHz (7mm) using VLBA with milli-arcsecond spatial resolution indicates the first
direct evidence that hourly time scale flares are localized within the inner 30×70
Schwarzschild radii of Sgr A*. We also show several cross correlation plots between
near-IR, millimeter and submillimeter light curves that collectively demonstrate the
presence of time delays between the peaks of emission up to three hours. The ev-
idence for time delays at millimeter and submillimeter wavelengths are consistent
with the source of emission being optically thick initially followed by a transition to
an optically thin regime. In particular, there is an intriguing correlation between the
optically thin near-IR and X-ray flare and optically thick radio flare at 43 GHz that
occurred on 2007 April 4. This would be the first evidence of a radio flare emission
at 43 GHz delayed with respect to the near-IR and X-ray flare emission. The time
delay measurements support the expansion of hot self-absorbed synchrotron plasma
blob and weaken the hot spot model of flare emission. In addition, a simultaneous
fit to 43 and 84 GHz light curves, using an adiabatic expansion model of hot plasma,
appears to support a power law rather than a relativistic Maxwellian distribution of
Subject headings: accretion, accretion disks — black hole physics — Galaxy: center
– 2 –
The black hole at the center of our own galaxy was first detected as the radio source Sgr A*
over 30 years ago (Balick & Brown 1974). It was found to lie at the center of a cluster of young
massive stars. Submillimeter and far-infrared observations showed that Sgr A* is encircled by a
torus of gas approximately 10 light-years across, which orbits with a speed of 100 km s−1(e.g.
Genzel & Townes 1987). The gravity required to hold onto this material implies a mass of several
million solar masses, although a portion of this is contributed by the stars in the stellar cluster.
These measurements suggested that Sgr A* could be a black hole.
through studies of the light distribution of stars in the cluster, as well as the motions of ionized and
molecular gas clouds orbiting Sgr A*. These measurements implied a mass of approximately 3–4
million times that of the sun (Genzel 2000; Genzel & Townes 1987) lying within a third of a light
year of the radio source. Recently, more precise measurements of fast moving stars in close orbits
around Sgr A* have conclusively demonstrated that it has a mass of ∼ 4×106M⊙(Ghez et al. 2005;
Eisenhauer et al. 2005; Ghez et al. 2008; Sch¨ odel et al. 2002; Gillessen et al. 2009) and that the size
of the radio source is about ∼4 times its Schwarzschild radius at 230 GHz (Rs) (Doeleman et al.
2008). This dark, massive object has also been uniquely identified through the proper motion of
the radio source, which show that Sgr A* must contain > 4 × 105M⊙(Reid & Brunthaler 2004).
Taken together, these measurements provide strong evidence that Sgr A* is a black hole with mass
∼ 4×106M⊙. No other known category of astrophysical object can easily fit so much mass into a
More detail was provided
1Dept. of Physics and Astronomy, Northwestern University, Evanston, Il. 60208
2STScI, 3700 San Martin Drive, Baltimore, MD 21218
3Department of Physics and Engineering, Macquarie University, Sydney NSW 2109, Australia
4Dept. of Physics, University of Alberta, Room #238 CEB, 11322-89 Avenue, Edmonton AB T6G 2G7, Canada
5Adler Planetarium and Astronomy Museum, 1300 South Lake Shore Drive, Chicago, IL 60605
6Cal Tech, Jet Propulsion Laboratory, Pasadena, CA 91109
7Max-Planck-Institut f¨ ur Radioastronomie, Auf dem Huegel 69, 53121 Bonn, Germany
8Harvard-Smithsonian CfA, 60 Garden Street, Cambridge, MA 02138
9Oberlin College, Dept. of Physics and Astronomy, Professor 110 N. St.,Oberlin, OH 44074
10National Radio Astronomy Observatory; University of Chicago, 5640 South Ellis Avenue, Chicago IL 60637
11Observatoire astronomique de Strasbourg, Universit´ e de Strasbourg, NRS, INSU, 11 rue de l’Universit´ e, 67000
12Max-Plank-Institut f¨ ur Extraterrestrische Physik 1312, D-85471, Garching, Germany
13Radio Astronomy Lab, 601 Campbell Hall, University of California, Berkeley, CA 94720
14Institut de RadioAstronomie Millimetrique, 300 rue de la Piscine, Domaine Universitaire 38406 Saint Martin
d’Heres, France, on leave to IRAM Granada, Spain
15Mizusawa VLBI Observatory, National Astronomical Observatory of Japan, Mizusawa, Oshu, Iwate 023-0861,
16School of Physics, University of Western Australia, 35 Stirling Highway, Crawley, WA, 6009, Australia
17Service d’Astrophysique / IRFU / DSM, CEA Saclay, Bat.
AstroParticule & Cosmologie (APC) / Universit´ e Paris VII / CNRS / CEA / Observatoire de Paris Bat. Condorcet,
10, rue Alice Domon et L´ eonie Duquet, 75205 Paris Cedex 13, France
709, 91191, Gif-sur-Yvette Cedex, France and
– 3 –
This massive black hole is a hundred times closer to us than the next nearest example, presenting an
unparalleled opportunity to closely study the process by which gas is captured by black holes. It is
therefore the subject of intense scrutiny. The energy radiated by Sgr A* is thought to be liberated
from gas that is falling into the black hole after being captured from the powerful winds of members
of its neighboring cluster of massive stars (e.g., Melia 1992). The broad band spectrum of Sgr A*
peaks at submillimeter wavelengths (Zylka et al. 1992; Falcke et al. 1998); this is thought to be
the dividing line between optically thick and optically thin emission at low and high frequencies,
respectively. The bolometric luminosity of Sgr A* ∼ 100
below that predicted given its expected rate of capture of material from stellar winds, prompting
a number of theoretical models to explain its very low efficiency (Narayan et al. 1995; Liu & Melia
2001; Yuan et al. 2003; Goldston et al. 2005; Liu et al. 2004; Falcke et al. 2009).
L⊙ is several orders of magnitudes
Now that the quiescent spectrum of emission from Sgr A* has been characterized from radio to X-
rays, attention has turned to variability of emission in multiple wavelengths. These measurements
probe the structure and the physical parameters of the hot plasma in the vicinity of the black hole
by measuring the time variations of its flux in different wavelength bands as well as their cross-
correlation with each other. Flaring activity on < 1 − 4 hour time scale is seen in all wavelength
bands in which quiescent emission has been detected.
Flaring X-ray emission from Sgr A* has been detected and has been argued to originate within
a few Schwarzschild radii of the ∼ 4 × 106M⊙black hole (Baganoff et al. 2003; Goldwurm et al.
2003; Porquet et al. 2003; B´ elanger et al. 2005). At near-IR (NIR) wavelengths (Genzel et al. 2003;
Yuan et al. 2003; Ghez et al. 2004; Hornstein et al. 2007), flare emission from Sgr A* is shown to
be due to optically thin synchrotron emission, whereas the long-wavelength flaring activity in
submillimeter, millimeter and centimeter bands is due to optically thick synchrotron emission.
The exact frequency at which the transition from optically thick to thin flare emission occurs is
A variety of mechanisms have been proposed to explain the origin of the variability of Sgr A*. Many
of these models have considered different energy distributions for the relativistic particles to explain
the origin of submillimeter emission (Markoff et al. 2001; Yuan et al. 2002; Melia 2002; Liu & Melia
2002; Yuan et al. 2003; Nayakshin & Sunyaev 2003; Eckart et al. 2004, 2006a,b; Yusef-Zadeh et al.
2006a; Gillessen et al. 2006; Goldston et al. 2005; Liu et al. 2006; Falcke et al. 2009); Melia and
Falcke (2001 and references therein). The direction that has been taken in the past in interpreting
the flaring activity of Sgr A* is within one of the established paradigms for the accretion flow
that have been developed based on the time-averaged emission – for example a thin accretion
disk, a disk and jet, outflow, an advection-dominated accretion flow, radiatively inefficient accre-
tion flow, accretion disk inflow/outflow solutions (Melia 1992; Yuan et al. 2003; Falcke & Markoff
2000; Falcke et al. 2009; Narayan et al. 1998; Blandford & Begelman 1999) and then the predicted
spectrum is compared with the observed spectrum.
We have recently analyzed the NIR flaring of Sgr A*, which is produced by synchrotron emission
from a transient population of particles produced within ∼ 10 Schwarzschild radii of the massive
black hole (Genzel et al. 2003; Eckart et al. 2006a; Gillessen et al. 2006).
∼ 2−3 hour duration of submillimeter flares could not be due to synchrotron cooling when observed
simultaneously with a NIR flare (estimated to be ∼ 20 minutes and ∼ 12 hours at 1.6µm [188 THz]
and 850µm [350 GHz], respectively). The decline in submillimeter light curves was interpreted to
be due to adiabatic cooling associated with expansion of the emitting plasma (Yusef-Zadeh et al.
2006a,b) under the assumption that the same accelerated population of particles is responsible
We argued that the
– 4 –
for NIR and submillimeter emission. Time delays detected between peaks of flare emission at
radio, submillimeter and NIR/X-rays wavelengths are consistent with this picture (e.g., Yusef-
Zadeh et al. 2006b; Marrone et al. 2008; Yusef-Zadeh et al. 2008; Meyer et al. 2008; Eckart
et al. 2008). However, the lack of long simultaneous coverage have not placed strong constraints
in time delay measurements, especially between radio and NIR wavelengths. Simple modeling of
the total and polarized intensity of the hot expanding plasma provide predictions that can be
tested observationally by carrying out observational campaigns such as the one we coordinated
during April 2007 to examine the mechanisms for the variability, with implications on the nature
of the accretion flow. The results presented here are the third in a series of papers that came
from the multi-wavelength observing campaign that took place on 2009 April (Porquet et al. 2008;
Dodds-Eden et al. 2009). The results of soft γ-ray observations are given separately (Trap et al.
The structure of this paper is as follows. §3 presents light curves of all the useful data that were taken
in this campaign, following the observational details described in §2. In §4 we analyze the statistical
properties of flare emission and the corresponding spectral and power spectrum distributions at
NIR wavelengths, as well as cross-correlation analysis of light curves. We then discuss in §5 the
origin of X-ray emitting flares and provide observational support for the expanding hot plasma
model of flare emission. The polarization results will be given elsewhere.
2. Observations and Data Reduction
The primary purpose of observations made during 2007 April 1-11 was the coordination of several
telescopes operating at many wavelengths to monitor the emission from Sgr A* and measure the
time evolution of its spectrum. There were a total of 13 observatories that participated in this cam-
paign, including XMM-Newton, the Hubble Space Telescope (HST), the International Gamma-Ray
Astrophysics Laboratory INTEGRAL, the Very Large Array (VLA) of the National Radio Astron-
omy Observatory18(NRAO), the Very Long Baseline Array (VLBA18), the Caltech Submillimeter
Observatory (CSO), the Very Large Telescope (VLT), the Submillimeter Array (Blundell 2004),
the 30m Pico Veleta Telescope of the Institute for Millimeter Radioastronomy (IRAM), the Sub-
millimeter Telescope (SMT), the Nobeyama Millimeter Array (NMA), the Combined Array for
Research in Millimeter-wave Astronomy (CARMA), and the Giant Meterwave Radio Telescope
(GMRT). The campaign was organized by first obtaining observing time with XMM-Newton (PI:
D. Porquet) and HST (PI: F. Yusef-Zadeh), and then coordinating the ground-based facilities to
the allotted space-based schedules.
Figure 1 shows the schedule of all observations and their rough durations.
Newton and VLT/NACO observations have already been reported in Porquet et al. (2008), and
(Dodds-Eden et al. 2009), respectively. Summary of the results from VLT/VISIR and INTEGRAL
observations has also been given by Trap et al. (2009). Briefly, XMM observations were carried out
using three observations for a total of 230 ks blocks of time during 2007 March 30 to April 4, and
the VLT/NACO observations took place between 2007 April 1–6 using H (1.66 µm), K (2.12µm),
and L’ (3.8µm) bands. INTEGRAL observations took place in parallel to XMM-Newton on April
1 and 4 for a total effective exposure time of 212 ks for IBIS/ISGRI (20–100 keV) and 46 ks for
Details of XMM-
18The National Radio Astronomy Observatory is a facility of the National Science Foundation, operated under a
cooperative agreement by Associated Universities, Inc.
– 5 –
JEM-X 1 (3–20 keV).
2.1. HST NICMOS: 1.45µm and 1.70µm
We obtained 40 orbits of HST observations using the NICMOS camera 1, with the orbits distributed
over seven consecutive days between 2007 April 1–7. The observations used the NICMOS F145M
and F170M filters, with exposure times of 144 sec in each filter and readout samplings of ∼16 sec
within each exposure. Use of this pair of filters has several advantages. First, they have no overlap
in wavelength space and are therefore suitable for spectral index measurements. Second, they are
well matched to one another in terms of throughput, so that we can use identical exposure and
readout times, thus producing time series data that are evenly sampled in both filters. This has
great advantages for making periodicity measurements. Third, their relatively high throughput
also allows us to use relatively short exposure times, so that we can cycle back and forth between
the two filters fairly rapidly. The near-simultaneous observations then allow us to make meaningful
spectral index measurements of flare events. This is especially true for the peaks of flare events,
where the flux from Sgr A* does not change rapidly over the course of several minutes or more.
The spectral index distribution can not be measured accurately, however, during the rise and fall of
flare events because the overall flux is changing more rapidly than the cadence of our filter cycling
The IRAF “apphot” routines were used to perform aperture photometry of sources in the NICMOS
Sgr A* field, including Sgr A* itself. For stellar sources the measurement aperture was positioned on
each source using an automatic centroiding routine. This approach could not be used for measuring
Sgr A* because its signal is spatially overlapped by that of the orbiting star S0-2 and S17. Therefore
the photometry aperture for Sgr A* was positioned by using a constant offset from the measured
location of S0-2 in each image. The offset between S0-2 and Sgr A* was derived from the orbital
parameters given by Ghez et al. (2003). The position of Sgr A* was estimated to be 0.16′′south
and 0.01′′west of S0-2 at the time of the HST observations. To confirm the accuracy of the position
of Sgr A*, two images of Sgr A* taken before and during a flare event were aligned and subtracted,
which resulted in an image showing the location of the flare emission.
We used a measurement aperture with a diameter of 3 detector pixels, which corresponds to ∼ 0.13′′.
This size was chosen as a suitable compromise between wanting to maximize the fraction of the
Sgr A* PSF included in the aperture, while at the same time limiting the amount of signal coming
from the wings of the PSF from the adjacent star S0-2. We derived aperture correction factors by
making measurements of a reasonably well isolated star in the field through a series of apertures of
increasing size. The aperture corrections, which convert the fluxes measured through our 3-pixel
diameter aperture to a semi-infinite aperture, are 2.91 and 3.40 for the 1.45 and 1.70 µm bands,
respectively. Absolute calibration of the photometric measurements was accomplished using the
latest calibrations for the NICMOS F145M and F170M obtained from STScI.
De-reddened fluxes were computed using the appropriate extinction law for the Galactic center
(Moneti et al. 2001) and their extinction value of A(K)=3.3 mag. These translate to extinction
values for the NICMOS filter bands of A(F170M)=5.03 mag and A(F145M)=6.52 mag, which then
correspond to correction factors of 103 and 406, respectively. Because there is some contribution
from the neighboring star S0-2 within our measurement aperture, we have determined the flux of
Sgr A* when it is flaring by subtracting the mean flux level measured during “quiescent” episodes.
The resulting net flux can therefore be safely attributed to Sgr A* flares.
– 6 –
2.2.VLA: 43 and 22 GHz
We used a fast-switching technique to observe Sgr A* simultaneously using the VLA D configuration
at 43.3 GHz (7mm) and 22.4 GHz (13mm) GHz. These observations took place on 2007 April 1–4,
each lasting for ∼7 hours, using 8 VLA and 18 eVLA antennas. The two IFs were separated by
50MHz in each observation, except for those on 2007 April 4, when the two IFs were centered at
non-standard frequencies of 43.1851 and 43.5351 GHz, which corresponds to a frequency separation
of 350 MHz. The separation was used to carry out polarization measurements, the result of which
will be given elsewhere.
We cycled between Sgr A* and the fast-switching calibrator 17444-31166 (2.3 degrees away from
Sgr A*) for 90 sec and 30 sec, respectively, throughout the observation. On 2007 April 4, we also
used the fast-switching phase calibrator 17459-28204, which is weaker, but closer (∼ 21′) to Sgr A*.
3C286 was used as the flux calibrator and NRAO530 was observed as a polarization and additional
phase calibrator. The light curves at 43 GHz restricted data to a uv range greater than 90 kλ with
full width at half point of 2.45′′×1.3′′(PA=−40). We used NRAO530 for pointing every 30 minutes;
the bootstrapped flux of NRAO530 at 43 GHz is 2.43±0.05 Jy. At 22 GHz, the strong continuum
emission from ionized gas associated with extended features surrounding Sgr A* overwhelmed the
flux of Sgr A* itself, making the variability analysis uncertain. Therefore the 22 GHz data are not
useful and are not presented here. In all the measurements presented here, we used only antennas
that had constant gain curves with similar values, thus many of the eVLA antennas were not used.
In all cases, at least two and sometimes three phase calibrators were used in order to ensure that
amplitude variability or calibration errors of one of the calibrators would not be introduced into the
light curve of Sgr A*. In the case of multiple phase calibrators, the same calibrator used to calibrate
the gains of Sgr A* was used to cross-calibrate the other calibrators. In cases where calibrator light
curves are shown as alongside those of Sgr A*, they were obtained from cross-calibration using
one of the other phase calibrators and not from self-calibration. Additionally, as a check, all light
curves of Sgr A* made using phase calibrations from the principle phase calibrator (usually 17444-
31166) were compared against light curves of Sgr A* made using the other phase calibrators. These
comparisons were used to identify bad data in the calibrators and after editing and recalibration
light curves using different calibrators were consistent.
For the final light curves of Sgr A*, the data were calibrated using the principle phase calibrator
(17444-31166). A phase self-calibration was applied to Sgr A* before the determination of a light
curve. No amplitude self-calibration was done to Sgr A* or any backup phase calibrators, whose
light curves are shown as a reference, since amplitude calibration would remove time variation from
the light curves. After phase self-calibration, large images were made and found no confusing point
sources above the rms noise (typically 2.5 mJy beam−1in a full run at 43 GHz) when the selected
uv data were greater than 90kλ.
In order to derive the light curve in the visibility plane, the Astronomical Image Processing System
(AIPS) task DFTPL was used. DFTPL plots the direct Fourier transform of a vector averaged
set of measured visibilities as a function of time. Since we use this on data that has been phase
self-calibrated using the point source Sgr A*, the vector average gives the flux of Sgr A*. Visibilities
are averaged in bins with defined time widths, however, since the number of visibilities in each bin
varies, the error associated with the average will not be constant and are derived at each time
– 7 –
2.3.VLBA: 43, 22, and 15 GHz
We observed Sgr A* with the VLBA in two different experiments. One took place on 2007 April 1,
5 and 11 at 43 GHz under program BR124. All observations employed four 8 MHz bands in dual
circular polarization each. These observations were made at 43 GHz and involved rapid switching
between Sgr A* and the two background continuum sources J1748-291 and J1745-283. Sources
were changed every 15 seconds in the sequence Sgr A* – J1748-291 – Sgr A* – J1745-283 – Sgr A*,
yielding an on-source time of ∼ 10 seconds. Before, in the middle, and after each observation
16 quasars were observed within ∼ 40 minutes. NRAO530 was also observed as fringe-finder.
The total observing time including the quasars was 8 hours for each observation. The data were
correlated with 16 spectral channels per frequency band and an integration time of 0.131 seconds.
The ∼ 1 hour gaps in the light curves of Sgr A* are due to geodetic measurements.
In the second experiment (proposal code BB230), the observations on 2007 April 2 and 10 involved
rapid switching between two frequencies on Sgr A* . We observed 43 and 22 GHz on 2007 April 2
and 43 and 14 GHz on April 10. We changed the receiver every 20 seconds, yielding an on-source
time of ∼10 seconds for each frequency. The observations were interrupted three times by 20 minute
observations of four different quasars, including the fringe-finder 3C345. All data were correlated
with 16 spectral channels per band and an integration time of one second. The total observing
time including the quasars was 6 hours for each observation.
The VLBI data were edited and calibrated using standard techniques in AIPS. First, we applied
the latest values of the Earth’s orientation parameters. A-priori amplitude calibration was applied
using system temperature measurements and standard gain curves. We performed a “manual phase-
calibration” using the data from NRAO530 or 3C345 to remove instrumental phase offsets among
the three frequency bands. Then, we fringe fitted the data from Sgr A* using only the five inner
VLBA antennas (PT, KP, FD, OV, and LA). Then, we discarded all data with elevations below 15◦
and performed one round of phase self-calibration on Sgr A*. Finally, we divided the calibrated
uv-data by a model described in Bower et al. (2004), i.e. an elliptical Gaussian component of
0.71×0.41 mas with a position angle of 78◦. Lightcurves were extracted form the uv-data with the
AIPS task DFTPL.
2.4. CSO: 350µm, 450µm & 850µm
Nightly observations of Sgr A* were made at three wavelengths over the period 2007 April 1–6 UT,
using the SHARC-II camera. The observations on April 1–5 were made with the SHARP imaging
polarimeter module installed (Li et al. 2008), using the 350 µm half-wave plate. The observing
bands (selected by a cryogenic filter) were 450 µm on April 1–3 and 350 µm on April 4–5. The
observations on April 6 were made at 850 µm with SHARP removed from the optical path.
In this paper, we report only the total intensity measurements.
polarization instrument, and because the observations were made over cycles of half-wave plate
angles that fully modulate the polarization, our total intensity results are insensitive to the polar-
ization of the source.
Because SHARP is a dual-
Except for the stepping of the polarimeter half-wave plate between integrations, the observing and
analysis method was similar to past CSO observations (Yusef-Zadeh et al. 2006a, 2008). We used
Lissajous scans with typical full amplitudes of 100′′in both azimuth and elevation. The instanta-
neous field of view is 57′′× 57′′in polarimeter mode and is 154′′× 58′′without the polarimeter.
– 8 –
The measured beam sizes were 8.4′′at 350 µm, 10.1′′at 450 µm, and 18.8′′at 850 µm. We used
the Dish Surface Optimization System (DSOS) on April 1–5. Three quadrants were working fully
during the run, but the fourth quadrant of the system was available for only part of the run. For
observations at the elevation of Sgr A*, we expect no significant effect on the results from the non-
operational quadrant. Any change in the beam FWHM due to the status of the DSOS quadrant
was less than 3%.
Mauna Kea weather conditions were good overall during the April 1–6 campaign. Occasional thin
cirrus was observed visually and on satellite photos on April 3, 4, and 6; otherwise skies were clear.
Local humidity was < 30% during the observations. The wind speed for April 1 was noticeably
high (roughly 30 mph), but less than 20 mph on the other nights. The zenith atmospheric opacity
at 225 GHz was marginal for observations at 450 µm on April 1 (τ225≈ 0.07), as well as at 850 µm
on April 6 (τ225≈ 0.15), and rising in both cases. Atmospheric opacity on April 2–5 was excellent
and relatively steady, ranging from 0.03 to 0.06 at 225 GHz.
In producing the light curves for this paper, we reconsidered the image registration and absolute
calibration for all of our CSO observations of Sgr A* from 2004 September through 2008 May. The
absolute pointing of the images is based on hourly measurements of point-like calibration sources,
such as planets, and the pointing model for the telescope. This procedure appears to average
down in a reasonable manner. The 350 µm and 450 µm position that we measure for the variable
component of Sgr A* is within 0.3′′of the nominal position of α2000= 17h: 45m: 40s.03,δ2000=
−290: 00′: 28′′.1. The agreement at 850 µm, at which the telescope beam size is larger, is somewhat
worse at 0.9′′.
At any particular point in time, the telescope pointing model has only ∼ 2′′accuracy. Therefore,
we shifted the individual observations to align with the average of all the observations, using the
bright dust emission in the images as the reference. Subsequent photometry of Sgr A* assumes a
fixed position and beam size.
Minor changes have been made to the absolute calibration scale factor and Sgr A* “zero point”, in-
cluding data which have been published in the past (Yusef-Zadeh et al. 2006a, 2008; Marrone et al.
2008). For the scale factor, we have adopted the following brightness temperatures for calibration
at 350, 450, and 850 µm, respectively: Callisto (128, 122, 120 K), Neptune (61, 66, 81 K), and
Uranus (64, 70, 86 K), arranged in decreasing order of usage and with an estimated 10% uncertainty.
These brightness temperatures are not significantly different from our past assumptions. The zero
point relates to the difficulty of measuring the total flux of Sgr A* with ∼ 10′′resolution because
of confusion from surrounding dust emission. We estimate an additive uncertainty of 1 Jy in our
measurements of the absolute flux at 350 µm and 850 µm, and an additive uncertainty of 0.5 Jy
at 450 µm. To be consistent with the results published in this paper, the August–September 2004
measurements at 850µm reported by Yusef-Zadeh et al. (2006a) should be shifted upwards by ∼0.2
Jy; the 450 µm measurements for the same period are essentially unchanged. The 850 µm mea-
surements for July 2006, reported by Yusef-Zadeh et al. (2008) and (Marrone et al. 2008), should
be shifted upwards by ∼0.5 Jy. The 350 µm results for July 2005, reported by (Marrone et al.
2008), should be shifted upwards by ∼0.4 Jy; the 450 µm and 850 µm results for the same period
are essentially unchanged.
– 9 –
2.5. SMA: 230 GHz
The SMA observed Sgr A* on the nights of 1, 3, 4, and 5 April 2007, typically covering the interval
1200–1830 UT. On the first three nights the array was tuned to observe 231.9 (221.9) GHz in the
upper (lower) sideband, while on the last night the frequency was tuned to 246.0 (241.0) GHz. The
array was in its “compact-north” configuration, resulting in angular resolution of approximately
3′′. All eight antennas were used except on April 4, when one was lost to an instrument problem.
The SMA polarimetry system (Marrone & Rao 2008) was used in these observations to convert the
linearly polarized SMA feeds to circular polarization sensitivity, which prevents confusion between
linear polarization and total intensity variations.
The data were gain calibrated using the quasar J1733−130, which was observed approximately
every 10 minutes. The absolute flux density scale was derived from observations of Callisto and
has an uncertainty of 10%. To remove the effects of the extended emission that surrounds Sgr A*,
only projected baselines longer than 20 kλ were used in the light curve determination. Flux density
measurements were made by applying the quasar gains to the Sgr A* data, removing the average
phase on Sgr A* in each light curve interval via phase self-calibration, which reduces the effect
of baseline errors and phase drifts on the measurement, and fitting a central point source to the
calibrated visibilities. Errors in the flux density account for thermal noise, as well as the time-
variable uncertainty in the gain, which is estimated from the data themselves.
2.6. IRAM-30m Telescope: 240 GHz
Observations with the IRAM-30m telescope at Pico Veleta, Spain, were carried out on 2007 April 1-
4. Because of the low elevation of Sgr A* at Pico Veleta, gain drifts due to atmospheric fluctuations
are the most severe limitation to accurate flux monitoring. For the same reason, accurate peak-up
is important if flux variations are to be measured that are small with respect to the quiescent
flux. To account for both requirements, we alternated between Sgr B2 and Sgr A* with the
following procedure, applying a wobbling secondary mirror to remove the 240 GHz emission from
the atmospheric and from extended (> 70”) source structure. First, we pointed at Sgr B2 and
measured the position of its point-source component by fitting simultaneously a Gaussian and a
linear baseline (for a refined removal of extended emission) to the pointing subscans taken in on-
the-fly mode (two in azimuth direction, two in elevation). The positional correction was entered
and the procedure repeated to recover the correct flux. We used either the azimuth or elevation
subscan, depending on where the flux was larger (and thus a better peak-up was provided). Then
the antenna was moved to Sgr A*, where the same procedure was repeated. Thus, for each time
sample, there are two data points representing the flux of Sgr A*, one from the pointing, and
another one from the peaked up pointing. Both results were used if the pointing correction was
sufficiently small. Error estimates were made by comparing the results of subscans in the same
direction. To avoid effects due to instrumental and atmospheric gains drifts, only scaled fluxes
of Sgr A* were retained for further analysis, using Sgr B2 as a non-variable flux reference. Data
reduction was done with the MOPSIC software package19The average Sgr B2 flux density is
estimated to be 38.0±1.2 Jy and was derived using the HII region G10.62-0.38 as absolute flux
reference. The beam FWHM is 11′′2.
– 10 –
2.7.SMT: 250 GHz
Observations were undertaken at the SMT located at 3200m altitude on Mount Graham in eastern
Arizona Baars et al. (1999) using the 250 GHz channel of the facility’s four color bolometer Kreysa
(1990). This bolometer was used to observe at 250 GHz with a broad band, ranging between 200
and 290 GHz on 2007 April 1–4. We made use of the telescope’s beam switching mode, chopping
horizontally ±2′with the subreflector at a rate of 2 Hz along with an “off-on-on-off” observing
mode that shifted the position of the telescope every 10 seconds to remove any asymmetries in
the observations due to the chopping. Jupiter, Saturn, and Mars were used for focus and pointing
references confirming the telescope’s typical half power beam width at these frequencies of 30′′and
pointing accuracy of 2′′. While Jupiter was used to set the gain of the bolometer and skydips to
find the atmospheric opacity, NRAO 530, 1757-240 and G34.3 were also observed throughout the
observations as secondary calibrators to check the stability and repeatability of the measurements.
Finally, after splitting the data into 80 second increments (consisting of two iterations of the 40
second long “off-on-on-off” observing mode), the raw data were reduced using a version of the
standard GILDAS NIC reduction program customized for the four color bolometer. Because a
single calibrator was not used continuously during the first two days of observations, the flux
variation of Sgr A* was uncertain and thus the data are not presented here.
2.8. NMA: 150 & 230 GHz
Interferometric NMA observations were carried out simultaneously at 90 and 102 GHz in the 3-mm
band, and simultaneously at 134 and 146 GHz in the 2-mm band with bandwidth of 1024 MHz
on 2007 April 1–4. The 2 and 3 mm flux densities are measured to be 1.8±0.4 and 2.0±0.3 Jy,
respectively. The light curves of data from April 3 and 4 are presented using five and six antennas,
respectively. Th weather was bad on April 1 and 2 so we discarded the data on the first days of
observations. 1744-31 (J2000) was used as the phase calibrator and the data was binned every 3-4
minutes. The flux measurements of Sgr A* were estimated by fitting a point source model in the
uv plane restricted to distances > 20 kλ, in order to suppress the contamination from extended
components surrounding Sgr A*. The FWHM of the synthesized beam in the 2mm observation on
2007, April 4 is 6′′× 1.4′′. We used 3C279 as a passband calibrator and Neptune as the primary
2.9. CARMA: 94 GHz
Interferometric CARMA observations were done to observe Sgr A* on 2007, April 2-5. Observations
were made at 94 GHz using nine 6m diameter BIMA and six 10m diameter OVRO antennas with
the exception of observations on 2007, April 3 which did not include any OVRO antennas. In all
days, Uranus was used as the primary flux calibrator, 1744-312 as the complex gain calibrator and
1751+096 was used as a passband calibrator. The weather was poor for observing at 94 GHz on
the first half of 2007, April 3 and the second half of 2007 April 4. We did not include the data
during these times. Five frequency windows, each 469 MHz wide, were used at frequencies from 94
to 100 GHz. We used NRAO530 (1730-130) to cross calibrate 1744-312, in order to independently
track the amplitude stability of 1744-312. All calibration was done using MIRIAD package and
calibrated visibility data for each day were read into AIPS and the DFTPL task was used to extract
– 11 –
light curves for the source.
2.10. GMRT: 1.28 GHz
We observed Sgr A* using Giant Meterwave Radio Telescope (GMRT) in 1280, 610 and 325 MHz
frequencies with central observation time on MJD 54195.0, 54191.1 and 54190.1 (5.0 April, 1.1
April and 31.1 March 2007 UT) respectively. GMRT20consists of thirty fully steerable parabolic
antenna array, where fourteen antennas are randomly distributed in 1 km area and rest of the
sixteen antennas are placed in three arms, spread over 25 km area, forming nearly a shape like ‘Y’.
The diameter of each antenna is forty five meter. Observation band-width in each frequency was
32 MHz and integration time was 16.9 second. The source was observed for 6.1, 4.2 and 7.0 hr in
1280, 610 and 325 MHz respectively. We have done flux calibration using 3C286 and 3C48 and used
Baars et al. (1977) for setting flux density scale. J1830-360 was used as phase calibrator. The bad
data and radio frequency interferences (RFIs) are eliminated from the data set and the source is
self-calibrated. The original data has channel width of 125 KHz in the spectral line mode. To take
care of effect of the band width smearing in low frequency, we did not averaged all the channels
after calibration but averaged 32, 16 and 8 channels in 1280, 610 and 325 MHz respectively (forming
effective channel width of 4, 2 and 1 MHz in 1280, 610 and 325 MHz). The images are corrected
for the beam-shape. Because of the strong emission from the nonthermal emission surrounding Sgr
A*, the light curve of Sgr A* could have not been obtained reliably at these low frequencies 330,
630 and 1280 MHz. This is mainly due to the instantaneous elongated beam shape which contains
extended structures surrounding Sgr A*.
3. Light Curves: Individual Telescopes
The results of XMM and VLT observations in X-ray and NIR have already been presented elsewhere
(Porquet et al. 2008; Dodds-Eden et al. 2009). A detailed account of INTEGRAL observations are
given elsewhere (Trap et al. 2009). To present all the data that were taken during this campaign,
we include the XMM and VLT light curves again here and briefly review the results of these
observations that have already been published elsewhere.
3.1. NICMOS Photometric Measurements
Figure 2a shows the observed variability of Sgr A* in the NICMOS 1.45 and 1.70µm bands, where
there is good agreement between the two bands. The observed “quiescent” emission levels of Sgr
A* in the 1.45 and 1.70µm bands are ∼ 32 and ∼ 38.5 mJy, respectively, but some fraction of
this total signal is due to the neighboring star S0-2. During flare events, the emission is seen to
increase by anywhere from a few percent to 25% above these levels. In spite of the somewhat lower
signal-to-noise ratio for the 1.45µm data, due to the somewhat lower sensitivity of the NICMOS
detector and increased effects of extinction, the flare activity is still easily detected in this band.
In order to confirm that the observed variability of Sgr A* is not due to either instrumental or
data reduction effects, we have compared the Sgr A* light curves to that of the star S0-2 and to a
– 12 –
region of background emission with the NICMOS images, as shown in Figure 2b. The photometric
measurements for Sgr A* show obvious signs of variability in six of the seven windows of HST
observations, while the corresponding light curves of S0-2 and the background remain quite stable.
The panels of Figure 3a-e present detailed light curves of Sgr A* and, for comparison, S0-2 for
each of the seven HST observing windows. These plots show the time-ordered measurements in the
1.45 and 1.70µm bands, where we have now subtracted the mean “quiescent” flux level, leaving the
net variations in emission for both Sgr A* and S0-2. All light curves are aperture and extinction
corrected. Each observing window consists of 5 to 7 HST orbits, with each orbit covering ∼ 46
minutes. We have identified flaring activity in at least one orbit in each of the seven observing
windows. These activities are identified in the light curves with labels designating the day (1–7)
and the flare even within the day (A–C). A typical flare event lasts between 10 and 40 minutes.
The amplitudes and durations of the events are similar to what was found in our earlier HST
observations (Yusef-Zadeh et al. 2006a).
To examine the short time scale variability in more detail, the 1.70µm light curves are shown in
Figure 4 with a sampling of 64 seconds. The Sgr A* and S0-2 light curves are qualitatively similar
to those in Figure 3, except for the finer sampling and we show only the 1.70µm band because
the 1.45µm data do not have sufficient signal-to-noise in this shorter integration period. There are
16 identified flaring events, all of which are shown in 45min periods in two panels of eight flares.
One type of fast fluctuation that we have detected is generally associated with the rise or fall of
bright flares, or at the peaks of bright flare emission, as seen for the flares 1A, 2A, and 5A. Similar
minute time-scale variability has also been detected by Dodds-Eden et al. (2009). Another type of
fluctuation is the point-to-point variability seen during some of the quiescent phases of low-level of
activity, such as flares 1B, 1C, 2B, 3A, 4B, 6A, and 7A.
3.2. VLT NIR and Mid-IR Observations
The VLT observations used multiple bands to observe Sgr A* on 2007 April 1–7, using the two
instruments NACO (NIR) and VISIR (mid-IR). The results of these observations, which included
the identification of seven flaring events are discussed in detail by Dodds-Eden et al. (2009). The
brightest flare detected at 3.80µm coincides with the brightest X-ray flare on April 4. Figure 5
shows a composite light curve of VLT observations with labeled flares using 3.8µm, 2.12µm and
1.66µm NIR bands. No NIR spectral index measurements are available for the detected flares.
However, a 3σ upper limit of 57 mJy is placed at 11.88µm for the bright 3.8µm flare on April 4
with a peak flux density of ∼30 mJy (see also Trap et al. 2009). The brightest NIR flare detected in
this campaign consists of a cluster of overlapping flares that last for about two hours. The second
brightest flare detected by the VLT is identified as #6 in Figure 5. This flare precedes the bright
NICMOS flare 5A (April 5), as shown in Figure 3e. These flares are components of another period
of flaring activity lasting for about two hours.
3.3. X-ray Flaring Activity
The X-ray light curves between 2 and 10 keV with a time binning of 144s are shown in Figure 6.
A total of five flares were observed: one in 2007 April 2 (labeled #1) with a peak X-ray luminosity
L2−10keV = 3.3 × 1034erg s−1and four on 2007 April 4 (labeled #2, #3, #4, #5) with peak
– 13 –
L2−10keV=24.6, 6.1, 6.3, and 8.9×1034erg s−1, respectively (Porquet et al. 2008). For the first
time, within a time interval of roughly half a day, an enhanced incidence rate of X-ray flaring
was observed, with a bright flare (#2, with a duration of 2900 s) followed by three flares of more
moderate amplitude (#3, #4, #5, with durations of 300, 1300, and 800s respectively). An enhanced
rate of X-ray flares, although with lower amplitudes, was also reported in Belanger et al. (2005)
when one moderate and two weak flares were detected within a period of eight hours. These rates
of X-ray activity (Porquet et al. 2008; B´ elanger et al. 2005) are clearly higher than the typical duty
cycle of one X-ray flare a day (Baganoff 2003). The brightest event on 2007 April 4 represents the
second-brightest X-ray flare from Sgr A* after the X-ray flare with Γ =2.2 ±0.3 on 2002, October 3,
on record with a peak amplitude of about 100 times above the 2–10keV quiescent luminosity21This
bright X-ray flare exhibits similar light-curve shape (i.e.,nearly symmetrical), duration (∼3 ks) and
spectral characteristics to the very bright flare observed on 2002, October 3 with XMM-Newton
(Porquet et al. 2003). Its measured spectral parameters, assuming an absorbed power law model
including the effects dust scattering, are NH= 12.3+2.1
quoted errors are at the 90% confidence level. Therefore, the two brightest X-ray flares observed
so far from Sgr A* exhibited similar soft spectra Γ ∼ 2.2 − 2.3. The spectral parameter fits of the
sum of the three following moderate flares, while lower (NH= 8.8+4.4
are compatible within the error bars with those of the bright flares. However, fixing the column
density at the value found for the brightest flare (i.e. NH= 12.3 × 1022cm−2) leads to a larger
photon index value for the sum of these moderate flares, i.e. Γ = 2.1±0.4.
−1.8× 1022cm−2and Γ =2.3 ±0.3 where the
−3.2×1022cm−2and Γ =1.7+0.7
3.4. 43 GHz Time Variability: VLA
Figure 7a,b shows light curves measured during April 1–4 at 43 GHz with the VLA, using 87sec and
300 sec sampling, respectively. The light curve of the phase calibrator 17444-31165, which itself is
cross calibrated by NRAO530, is flat and is shown at the bottom of each panel in Figure 7a. Since
NRAO530 is not the primary calibrator, it provides a second check on instrumental stability and
that its light curve was flat also.
The light curves of Sgr A* show variations on a variety of time scales from as short as 30 min to
longer than five hours at 43 GHz. The fluctuations on time scales of several hours ∼ 5−6 hours can
be seen in Figure 7a,b. The slow flux variation over 5-6 hours could, in principle, result from the
contamination of the emission by an asymmetric distribution of extended structures surrounding
Sgr A* especially when a compact configuration of the VLA is used. However, the contamination
of flux by extended emission is minimal for uv data > 90kλ (or 2.3′′) and the variability on several
hour time scale is intrinsic to Sgr A*. Previous high resolution data taken with a wide configuration
of the VLA have also shown the presence of flux variation of Sgr A* on such time scales (Yusef-
Zadeh et al. 2006a,b and 2008). The contamination of extended emission was clearly seen at low
elevations in the uv data < 90kλ at 43 GHz, and our 22GHz data taken simultaneously with 43
GHz data on 2007 April 1-4 were useless for time variability analysis because of the limited uv
range (i.e., < 70kλ).
Most of the power of the 43 GHz fluctuations in four consecutive days of observations appears to
fall in a range between 30 minutes and few hours, as best shown all light curves of Figure 7b. For
21No detection was made using INTEGRAL in the 20–40 keV and 40–100 keV energy bands, leading to 3σ upper
limits of 2.63 and 2.60×1035ergs s−1, respectively (Trap et al. 2009).
– 14 –
example, fluctuations with ∼1h time scale are detected at a level of 200 mJy in the April 1 and
April 2 light curves centered near 13h and 11:15h UT, respectively. The light curve of April 4
shows largest flux variations at a level of ∼40% are seen to increase flux density from 1.1 Jy at 9h
UT to 1.6 Jy near 15h UT. Another interesting feature of the April 4 light curve is the presence of
multiple weak fluctuations at a level of 50 mJy on a time scale of ∼20-30 minutes. Figure 7c shows
the light curves of April 4 for simultaneous observations at frequencies of 43.1851 GHz and 43.5351
GHz with a 30sec sampling time. The frequency separation between these light curves is 345 MHz.
We note at least five 20–30 minute fluctuations that are seen in both light curves. A more detailed
account of the power spectrum analysis of the time variability of Sgr A* in radio wavelengths will
be given elsewhere.
3.5. 14, 22 and 43 GHz Time Variability: VLBA
Figure 8a shows 43 GHz light curves based on VLBA observations on April 1, 5 and 11, with a 60
sec sampling time. Figure 8b shows the light curves at 22 GHz and 43 GHz on April 2 whereas
Figure 8c shows the light curves at 15 GHz and 43 GHz on April 10. The flux density of Sgr A* show
variations on several hour time scales in these VLBA observations at multiple frequencies. These
light curves show the first measurements of the flux variation of Sgr A* on a VLBA (milli-arcsecond)
scale at several frequencies.
Fluctuations in phase coherence and amplitude errors could produce significant changes in flux on
short timescales. However, it is unlikely that calibration errors are similar at two frequencies, thus
the flux variation on ∼5-hour time scale (Fig. b,c) is intrinsic to Sgr A*.
184.108.40.206 GHz Light Curve: VLBA and VLA Comparison
Because VLBA and VLA measurements on April 1 and 2 are taken simultaneously at 43 GHz, we
compared the two light curves, as shown in Figure 9a,b with a 300 sec sampling time, respectively.
The comparison of the light curves examines directly the localization of flaring events at radio
wavelengths. The largest fluctuations in both VLA and VLBA light curves appear to agree with
each other. Peaks with hourly time scale durations occur in both light curves near 13h UT on
April 1, as seen in Figure 9a. Similarly, the slow decreasing trend in the flux of Sgr A* over few
hours is seen in the light curves of 2007 April 2 at 43 GHz using both the VLA and VLBA, as
shown in Figure 9b. The behavior of the light curves on hourly time scales measured with VLBA
provides the first direct evidence that flaring activity arises from the innermost region of Sgr A* on
milliarcsecond (mas) scales. The size of the flare emission is dominated by interstellar scattering.
The general agreement between the VLA and VLBA light curves imply that flaring region that has
been detected is unresolved with the VLA.
There are also discrepancies between the two light curves. One is the different values of “average
levels” of flux taken in the light curves measured with the VLA and VLBA. In all the measurements
shown in Figures 7a and 8a,b the average-level of VLBA flux appears to be lower than than that of
the VLA by ∼ 200 mJy. The second discrepancy is the flux variations do not agree with each other
on small time scales in VLA and VLBA light curves. The uncertainty in the absolute flux density
calibration of Sgr A* at 43 GHz using VLBA and VLA could easily explain the first discrepancy.
It is possible that the emission from Sgr A* could be contaminated by extended emission from the
– 15 –
surrounding medium, as measured with the VLA, even though we have selected data with uv >
100kλ. This could explain why the VLA and VLBA light curves do not agree with each other on
10-15 minute time scales. Lastly, it is possible that these discrepancies could be explained by a
core-halo structure of emission from Sgr A* in which the halo component is resolved out in VLBA
observations. Future simultaneous VLA observations using its most extended array configuration
and VLBA should be able to examine closely the reason for these discrepancies.
3.6. CSO 350µm, 450µm, 850µm Light Curves
Figure 10 shows the light curves at three submillimeter wavelengths. The data have been smoothed
to increase the signal-to-noise ratio with a sampling time of ∼6.5 minutes. As at radio wavelengths,
the flux of Sgr A* appears to be varying on hourly time scales. The largest Sgr A* increase is
detected at the beginning of the observation near 13:30 UT on 2007 April 1. These light curves
show evidence for hourly and intraday variability at 450µm, at a level of 14%. The mean daily flux
of Sgr A* at 450µm is ∼ 3 ± 0.25 Jy.
Figure 10b shows some of the first variability of Sgr A* at 350µm. The flux increase on 2007,
April 4 over 5 hours is about 50% of the initial flux of Sgr A*. This steady increase of flux density
over several hours is seen to continue at 90 GHz (see section 3.10). Figure 10c shows the light
curve at 850µm on 2007, April 6. unlike the other submillimeter light curves shown here, this
light curve appears to show time variability on a time scale of ∼ 10 minutes as seen near 13h:20m
UT. Such sharp variations, though with low signal-to-noise values, at 850µm at such a short time
scale resembles the recent light curve obtained with a different instrument (LABOCA of APEX)
(Eckart et al. 2008). The reality of such a short time scale variation needs to be confirmed.
3.7. SMA: 230 GHz Light Curves
Figure 11 shows the light curves taken from four days of observations with the SMA at 230 GHz.
The 2007 April 1 data shows an asymmetric profile indicating a duration of possibly ∼ 4 hours
considering that there is a gap between 16h and 17h:30m UT. Similar submillimeter characteristics
have been seen recently at 850µm (Yusef-Zadeh et al. 2008). The 2007 April 3 light curve shows
an emission peak near 14h UT, before a slow decay that lasts for about 4 hours. The light curve
obtained with SMT on the same day and at the same wavelength showed the rising part of the
light curve suggesting that the duration of the flare on this day could be as long as 8 hours. The
2007 April 4 light curve shows a typical profile of submillimeter flare emission, except for a dip in
the flux at a level of 100 mJy near 14h UT. The April 5 data shows a light curve with multiple
peaks as the light curve decays. The typical time scale for this variation is between ∼1-2 hours.
The overall percentage of flux variation during 6 hours of observations is between 10% and 30%.
3.8. SMT 250 GHz Light Curves
The SMT light curves of Sgr A* and calibrators (in blue) for 2007, April 1-4 are shown in Figure 12.
Because SMT and SMA observed Sgr A* at the same wavelength considering the broad bandwidth
of the SMT, we compared the SMT light curves with those of SMA on April 1, 3 and 4. An increase
in flux of ∼1 Jy in the rising part of the light curve is seen between 10h UT and 14h UT on 2007,