Issue |
A&A
Volume 511, February 2010
|
|
---|---|---|
Article Number | A52 | |
Number of page(s) | 13 | |
Section | Extragalactic astronomy | |
DOI | https://doi.org/10.1051/0004-6361/200913073 | |
Published online | 09 March 2010 |
PKS 2005-489 at VHE: four years of monitoring with HESS and simultaneous multi-wavelength observations
HESS Collaboration
- F. Acero15 - F. Aharonian1,13 - A. G. Akhperjanian2 - G. Anton16 - U. Barres de Almeida8,
- A. R. Bazer-Bachi3 - Y. Becherini12 - B. Behera14 - W. Benbow1,
- K. Bernlöhr1,5 - A. Bochow1 - C. Boisson6 - J. Bolmont19 - V. Borrel3 - J. Brucker16 - F. Brun19 - P. Brun7 - R. Bühler1 - T. Bulik29 - I. Büsching9 - T. Boutelier17 - P. M. Chadwick8 - A. Charbonnier19 - R. C. G. Chaves1 - A. Cheesebrough8 - L.-M. Chounet10 - A. C. Clapson1 - G. Coignet11 - L. Costamante1,30,
- M. Dalton5 - M. K. Daniel8 - I. D. Davids22,9 - B. Degrange10 - C. Deil1 - H. J. Dickinson8 - A. Djannati-Ataï12 - W. Domainko1 - L. O'C. Drury13 - F. Dubois11 - G. Dubus17 - J. Dyks24 - M. Dyrda28 - K. Egberts1 - P. Eger16
- P. Espigat12 - L. Fallon13 - C. Farnier15 - S. Fegan10 - F. Feinstein15 - A. Fiasson11 - A. Förster1 - G. Fontaine10 - M. Füßling5 - S. Gabici13 - Y. A. Gallant15 - L. Gérard12 - D. Gerbig21 - B. Giebels10 - J. F. Glicenstein7 - B. Glück16 - P. Goret7 - D. Göring16 - M. Hauser14 - S. Heinz16 - G. Heinzelmann4 - G. Henri17 - G. Hermann1 - J. A. Hinton25 - A. Hoffmann18 - W. Hofmann1 - P. Hofverberg1 - M. Holleran9 - S. Hoppe1 - D. Horns4 - A. Jacholkowska19 - O. C. de Jager9 - C. Jahn16 - I. Jung16 - K. Katarzynski27 - U. Katz16 - S. Kaufmann14 - M. Kerschhaggl5 - D. Khangulyan1 - B. Khélifi10 - D. Keogh8 - D. Klochkov18 - W. Kluzniak24 - T. Kneiske4 - Nu. Komin7 - K. Kosack1 - R. Kossakowski11 - G. Lamanna11 - J.-P. Lenain6 - T. Lohse5 - V. Marandon12 - O. Martineau-Huynh19 - A. Marcowith15 - J. Masbou11 - D. Maurin19 - T. J. L. McComb8 - M. C. Medina6 - J. Méhault15 - R. Moderski24 - E. Moulin7 - M. Naumann-Godo10 - M. de Naurois19 - D. Nedbal20 - D. Nekrassov1 - B. Nicholas26 - J. Niemiec28 - S. J. Nolan8 - S. Ohm1 - J.-F. Olive3 - E. de Oña Wilhelmi1 - K. J. Orford8 - M. Ostrowski23 - M. Panter1 - M. Paz Arribas5 - G. Pedaletti14 - G. Pelletier17 - P.-O. Petrucci17 - S. Pita12 - G. Pühlhofer18,14 - M. Punch12 - A. Quirrenbach14 - B. C. Raubenheimer9 - M. Raue1,30 - S. M. Rayner8 - M. Renaud12,1 - F. Rieger1,30 - J. Ripken 4 - L. Rob 20 - S. Rosier-Lees11 - G. Rowell26 - B. Rudak24 - C. B. Rulten8 - J. Ruppel21 - V. Sahakian2 - A. Santangelo18 - R. Schlickeiser21 - F. M. Schöck16 - U. Schwanke5 - S. Schwarzburg18 - S. Schwemmer14 - A. Shalchi21 - M. Sikora24 - J. L. Skilton 25 - H. Sol 6 -
.
Stawarz23 - R. Steenkamp22 - C. Stegmann16 - F. Stinzing16 - G. Superina10 - A. Szostek23,17 - P. H. Tam 14 - J.-P. Tavernet19 - R. Terrier12 - O. Tibolla1 - M. Tluczykont 4 - C. van Eldik1 - G. Vasileiadis15 - C. Venter9 - L. Venter6 - J. P. Vialle11 - P. Vincent19 - M. Vivier7 - H. J. Völk1 - F. Volpe1 - S. J. Wagner14 - M. Ward8 - A. A. Zdziarski24 - A. Zech6
1 - Max-Planck-Institut für Kernphysik, PO Box 103980, 69029
Heidelberg, Germany
2 -
Yerevan Physics Institute, 2 Alikhanian Brothers St., 375036 Yerevan,
Armenia
3 -
Centre d'Etude Spatiale des Rayonnements, CNRS/UPS, 9 Av. du Colonel Roche, BP
4346, 31029 Toulouse Cedex 4, France
4 -
Universität Hamburg, Institut für Experimentalphysik, Luruper Chaussee
149, 22761 Hamburg, Germany
5 -
Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15,
12489 Berlin, Germany
6 -
LUTH, Observatoire de Paris, CNRS, Université Paris Diderot, 5 Place Jules Janssen, 92190 Meudon,
France
7 -
IRFU/DSM/CEA, CE Saclay, 91191
Gif-sur-Yvette, Cedex, France
8 -
University of Durham, Department of Physics, South Road, Durham DH1 3LE,
UK
9 -
Unit for Space Physics, North-West University, Potchefstroom 2520,
South Africa
10 -
Laboratoire Leprince-Ringuet, Ecole Polytechnique, CNRS/IN2P3,
91128 Palaiseau, France
11 -
Laboratoire d'Annecy-le-Vieux de Physique des Particules,
Université de Savoie, CNRS/IN2P3, 74941 Annecy-le-Vieux,
France
12 -
Astroparticule et Cosmologie (APC), CNRS, Université Paris 7 Denis Diderot,
10 rue Alice Domon et Léonie Duquet, 75205 Paris Cedex 13, France
13 -
Dublin Institute for Advanced Studies, 5 Merrion Square, Dublin 2,
Ireland
14 -
Landessternwarte, Universität Heidelberg, Königstuhl, 69117 Heidelberg, Germany
15 -
Laboratoire de Physique Théorique et Astroparticules,
Université Montpellier 2, CNRS/IN2P3, CC 70, Place Eugène Bataillon, 34095
Montpellier Cedex 5, France
16 -
Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1,
91058 Erlangen, Germany
17 -
Laboratoire d'Astrophysique de Grenoble, INSU/CNRS, Université Joseph Fourier, BP
53, 38041 Grenoble Cedex 9, France
18 -
Institut für Astronomie und Astrophysik, Universität Tübingen,
Sand 1, 72076 Tübingen, Germany
19 -
LPNHE, Université Pierre et Marie Curie Paris 6, Université Denis Diderot
Paris 7, CNRS/IN2P3, 4 Place Jussieu, 75252, Paris Cedex 5, France
20 -
Charles University, Faculty of Mathematics and Physics, Institute of
Particle and Nuclear Physics, V Holesovickách 2, 180 00, Czechoslovakia
21 -
Institut für Theoretische Physik, Lehrstuhl IV: Weltraum und
Astrophysik,
Ruhr-Universität Bochum, 44780 Bochum, Germany
22 -
University of Namibia, Private Bag 13301, Windhoek, Namibia
23 -
Obserwatorium Astronomiczne, Uniwersytet Jagiellonski, ul. Orla 171,
30-244 Kraków, Poland
24 -
Nicolaus Copernicus Astronomical Center, ul. Bartycka 18, 00-716 Warsaw,
Poland
25 -
School of Physics & Astronomy, University of Leeds, Leeds LS2 9JT, UK
26 -
School of Chemistry & Physics,
University of Adelaide, Adelaide 5005, Australia
27 -
Torun Centre for Astronomy, Nicolaus Copernicus University, ul.
Gagarina 11, 87-100 Torun, Poland
28 -
Instytut Fizyki Jadrowej PAN, ul. Radzikowskiego 152, 31-342 Kraków,
Poland
29 -
Astronomical Observatory, The University of Warsaw, Al. Ujazdowskie
4, 00-478 Warsaw, Poland
30 -
European Associated Laboratory for Gamma-Ray Astronomy, jointly
supported by CNRS and MPG
Received 6 August 2009 / Accepted 10 November 2009
Abstract
Aims. Our aim is to study the very high energy (VHE; E>100 GeV) -ray emission from BL Lac objects and the evolution in time of their broad-band spectral energy distribution (SED).
Methods. VHE observations of the high-frequency peaked
BL Lac object PKS 2005-489 were made with the High Energy
Stereoscopic System (HESS) from 2004 through 2007. Three simultaneous
multi-wavelength campaigns at lower energies were performed during the
HESS data taking, consisting of several individual pointings with the XMM-Newton and RXTE satellites.
Results. A strong VHE signal, 17
total, from PKS 2005-489 was detected during the four years of
HESS observations (90.3 h live time). The integral flux above the
average analysis threshold of 400 GeV is
3%
of the flux observed from the Crab Nebula and varies weakly on time
scales from days to years. The average VHE spectrum measured from
300 GeV to
5 TeV is characterized by a power law with a photon index,
.
At X-ray energies the flux is observed to vary by more than an order of
magnitude between 2004 and 2005. Strong changes in the X-ray
spectrum (
)
are also observed, which appear to be mirrored in the VHE band.
Conclusions. The SED of PKS 2005-489, constructed for the
first time with contemporaneous data on both humps, shows significant
evolution. The large flux variations in the X-ray band, coupled with
weak or no variations in the VHE band and a similar spectral behavior,
suggest the emergence of a new, separate, harder emission component in
September 2005.
Key words: galaxies: active - BL Lacertae objects: individual: PKS 2005-489 - gamma rays: galaxies - X-rays: individuals: PKS 2005-489
1 Introduction
PKS 2005-489 is one of the brightest BL Lac objects, at all wavelengths, in the Southern Hemisphere. It was initially discovered as a strong radio source in the Parkes 2.7 GHz survey (Wall et al. 1975) and later identified as a BL Lac object (Wall et al. 1986). It belongs to the complete 1-Jy radio catalog (Stickel et al. 1991), and is one of the few extragalactic objects detected in the EUV band (Marshall et al. 1995). Its redshift, z=0.071 (Falomo et al. 1987), is determined from weak, narrow emission lines observed during a low optical state.
PKS 2005-489 is classified as a High-frequency peaked BL Lac object (HBL; Giommi & Padovani 1994),
because of its high X-ray-to-radio flux ratio (Sambruna et al. 1995)
and because its broad-band spectral energy distribution (SED) peaks
in the UV-soft X-ray band. As is typical of HBLs, the X-ray spectrum is dominated
by the synchrotron emission of high-energy electrons.
The second SED hump is expected to peak in the GeV-TeV -ray band,
and is commonly believed to be produced
by the same electrons
up-scattering via the inverse Compton mechanism
seed photons of lower energy.
In the X-ray band, PKS 2005-489 has been studied extensively,
showing an extreme flux and spectral variability.
Large flux variations with correlated spectral hardening of
the generally steep spectrum
(photon index
)
were observed during five EXOSAT observations (Sambruna et al. 1994; Giommi et al. 1990).
Two ROSAT observations in 1992 confirmed the EXOSAT results and similarly
show a soft spectrum (
;
Sambruna et al. 1995).
A harder spectrum (
from 2 to 10 keV) was observed
in September 1996, during BeppoSAX observations of a brighter X-ray state (Padovani et al. 2001).
In October-November 1998 PKS 2005-489 underwent a period of exceptional activity,
with several strong X-ray flares. RXTE monitoring observations (Perlman et al. 1999)
were performed during the entire epoch, and the 2-10 keV flux reached
3
erg cm-2 s-1,
approximately 30 times higher than the ROSAT values. These RXTE observations
yielded a detection up to 40 keV and showed variations in the photon index
between
and 2.8. An X-ray flare alert also triggered BeppoSAX observations
on November 1-2, 1998. The X-ray spectrum was measured between 0.1 to 200 keV
and was characterized by a curved shape and harder photon indices,
in both the soft (
)
and hard (
)
X-ray bands
(< and >2 keV; Tagliaferri et al. 2001).
More recently, observations with the Swift satellite
have generally shown PKS 2005-489
in a low-flux, steep-spectrum state (
,
Massaro et al. 2008).
Strong correlations between X-ray and -ray emission
have been observed in many HBL
(e.g. Fossati et al. 2008; Pian et al. 1998; Aharonian et al. 2009; Maraschi et al. 1999; Krawczynski et al. 2004),
and are typically expected in a synchrotron-Compton scenario.
Therefore the very bright flux and large flux/spectral variability
at X-ray energies make PKS 2005-489 one of the most promising targets
for observing a similar behavior in the
-ray domain.
In the VHE (E>100 GeV) band, PKS 2005-489 was detected neither during
observations made between 1993 and 2000 by either the CANGAROO or
Durham groups (Roberts et al. 1998; Chadwick et al. 2000; Nishijima 2002; Roberts et al. 1999),
nor by HESS with a partial array
during its commissioning in 2003.
In 2004, HESS discovered
VHE -ray emission from PKS 2005-489
with a significance of 6.7
,
at a flux of a few percent of the Crab Nebula (Aharonian et al. 2005).
The measured spectrum was soft (
).
In the MeV-GeV band,
PKS 2005-489 is one of the few HBL detected by EGRET.
However, the observed significance is marginal:
3.7
above 100 MeV (Lin et al. 1996) and 4.1
above 1 GeV (Lamb & Macomb 1997).
It was instead detected by Fermi with high significance (>10
),
during the first three months of operation (2008, Aug.-Oct., Abdo et al. 2009).
Because of the high potential for strong -ray activity,
PKS 2005-489 has been monitored at VHE by HESS every year since its detection in 2004.
Several campaigns of coordinated observations with the X-ray satellites XMM-Newton and RXTE
were also performed. These simultaneous observations, sampling
the same particle distribution through two different emission processes,
represent a powerful diagnostic tool for probing the conditions of the inner blazar jet,
especially during flaring events (Coppi & Aharonian 1999).
In this article the results of all HESS observations taken from 2004 through 2007 are presented, together with the results of the multi-wavelength observations. These campaigns characterize the SED of PKS 2005-489 during different states and, for the first time, sample both humps of the SED simultaneously. A re-analysis of the 2004 HESS data is also provided, which benefits from an improved calibration of the absolute energy scale with respect to the previously published result.
Table 1: Results from long-term HESS observations of PKS 2005-489.
2 HESS observations and analysis technique
PKS 2005-489 was observed with the HESS array (Hinton 2004)
for a total of 158.0 h
(352 runs of 28 min each) from 2004 through 2007. During these observations
the array tracked a position offset from the blazar by 0.5
in alternating directions to
enable both on-source observations and simultaneous estimation
of the background induced by charged cosmic rays.
A total of 207 runs pass the standard HESS data-quality
selection, yielding an exposure of 90.3 h live time at
a mean zenith angle
.
The results presented here were generated using
the standard HESS calibration methods (Aharonian et al. 2004)
and analysis tools (Benbow 2005), with the
standard cuts event-selection criteria
(except for the 2007 spectrum, see Sect. 3.2).
On-source data were taken from a circular region
of radius
centered
on PKS 2005-489, and the background (off-source data) was estimated using
the Reflected-Region method (Berge et al. 2007).
Equation (17) in Li & Ma (1983) was used to calculate
the significance of any excess. All VHE integral fluxes
reported throughout this article were calculated assuming
the time-average photon index of
determined
in Sect. 3.2.
The PKS 2005-489 observations presented here span
four years (2004-2007) of HESS data taking. During this time
the optical throughput of the instrument decreased,
because of the degradation of the reflective surfaces
of the mirrors and Winston cones, as
well as accumulation of dust on the optical elements.
For the entire data sample,
the optical efficiency has decreased by an average of
28% compared to a newly commissioned instrument,
with its mirrors installed in Oct. 2001, Dec. 2002, June 2003,
and August 2003 on CT3, CT2, CT4, and CT1, respectively.
To minimize the effects of long-term variation in
the optical efficiency of the HESS array, the estimated
energy of each event was corrected
using the ratio of efficiencies determined on a run-wise basis
from simulated and observed muons (Aharonian et al. 2006a). After accounting for
the decreasing optical throughput of the
HESS array, the average energy threshold of the
analysis at
is 400 GeV.
3 HESS results
![]() |
Figure 1:
Distribution of |
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PKS 2005-489 was clearly detected in each of the four years (2004-2007)
that it was observed by HESS. A total of 1233 excess events, corresponding
to a statistical significance of 16.7 standard deviations (),
was detected from the direction of the blazar. The results
of the HESS observations are given in Table 1, which shows
the dead time corrected observation time,
the number of on and off-source events, the on/off normalization (
),
the excess and the corresponding statistical significance, for various temporal
breakdowns of the HESS data sample.
Figure 1 shows the on-source and normalized off-source
distributions of the square of the angular difference between
the reconstructed shower position
and the source position (
)
for all observations.
As can be seen in the figure, the distribution
of the excess (i.e. the observed signal) is very similar
to what is expected from a simulated point-source of VHE
-rays
at comparable zenith angles. The off-source distribution
is approximately flat in
,
as expected.
The map of excess counts is well-fit by a two-dimensional
Gaussian with a centroid located at
and
.
As expected, the fit location of HESS J2009-488 is consistent
with the position
(
,
)
of the blazar (Johnston et al. 1995).
The upper limit (99% confidence level)
on the extension of HESS J2009-488 is 1.0'.
3.1 VHE flux
The observed integral flux above 400 GeV for the entire data set is
I(>400 GeV
cm-2 s-1. This corresponds to
2.9% of
the flux above 400 GeV from the Crab Nebula, as determined by HESS (Aharonian et al. 2006a).
Figures 2 and 3
show the flux measured for each dark period and night, respectively.
![]() |
Figure 2: Integral flux, I(>400 GeV), measured by HESS from PKS 2005-489 during each dark period of observations (i.e. moonless night time within a month). Only the statistical errors are shown. The horizontal line represents the average flux for all the HESS observations. The four horizontal line segments represent the average annual flux observed during the corresponding years (see Table 1). |
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The integral flux I(>400 GeV) observed
during various epochs is reported in Table 1,
together with the
and corresponding probability, P(
),
for fits of a constant to the data when binned by nights within each dark period, by dark periods within
a year, and by year within the total observations. There are weak indications
(P(
)
of variability on annual time scales, monthly time scales in 2006,
and nightly time scales in the dark periods of September 2005 and August 2006.
The combined variations on different time scales lead to an indication
of variability in the overall monthly curve
(
for 13 NDF, with P(
,
Fig. 2).
Although no variability is seen inside the other epochs (e.g. monthly time scales in 2004, 2005,
or 2007), variations in amplitude comparable to the statistical errors
cannot be ruled out.
3.2 VHE spectra
The time-average photon spectrum for the entire data set is shown
in Fig. 4. All points in the energy range
300 GeV to 5.3 TeV are significant except the
last (1.5
)
at
4.6 TeV. The data are well-fit,
of 9.9 for 8 degrees of freedom,
by a power law (dN/dE=I400 (E/400 GeV)
)
with a photon index
.
Removing the
4.6 TeV point does not significantly
affect the fit result. No evidence is found
for significant features, such as a cut-off or break, in the energy spectrum,
also considering the upper limits at higher energies.
The time-average spectra measured during each year of data-taking are shown
in Fig. 5. The results of the best
fits
of a power law to these data are shown in Table 2.
All spectra are generated with standard cuts (Benbow 2005),
except for 2007. Given the
short exposure, low flux, steep spectral slope, and degradation
of the optical efficiency, a spectrum for 2007 could only be generated
with the spectrum cuts (Aharonian et al. 2006b),
which lower the energy threshold of the analysis at the expense
of sensitivity at higher energies.
In Table 2, the epoch, lower and upper energy bounds, photon index, differential flux
normalization at 400 GeV (I400),
,
degrees of freedom (NDF), and
probability P(
)
for each fit are given.
The
probability for a fit of a constant
to the annual
values is 0.30.
![]() |
Figure 3: Integral flux, I(>400 GeV), measured by HESS from PKS 2005-489 during each night of observations. The individual plots represent each of the four years (2004-2007) of data taking. Only the statistical errors are shown. The horizontal lines represent the average annual flux observed during the respective year. The vertical lines at MJD 53282, 53640, and 53642 represent the nights of XMM-Newton observations. The epoch of the RXTE observations is between the vertical lines at MJD 53609 and 53623. |
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Table 2:
Best
fits of a power-law to the various
spectra of PKS 2005-489 measured by HESS.
3.3 Correction to 2004 HESS data
The data previously published (Aharonian et al. 2005)
for HESS observations of PKS 2005-489 in 2004 were
not corrected for a 19% average decrease
in the optical efficiency until 2004.
In addition, the early simulations neglected the small gaps between
Winston cone apertures, the combined effect being an overall correction
of 22% compared to the initial analysis.
Figure 6 illustrates
the effect of correcting the energy of individual events for the relative
optical efficiency of the system, as described earlier.
Comparing (see Table 2)
the spectrum for 2004 determined here to
the previously published one shows
significant differences in the flux normalization
(I400) but not in the photon index ().
The flux measured in 2004 is three times
higher than previously published
, because of the
steep spectrum of the source.
It should be noted that, according to the model for the decrease
in the optical efficiency of the HESS system, the simulated values
match the actual values measured in 2003, therefore the upper limit
published for 2003 HESS observations is unchanged. Extrapolating
the 2003 upper limit
(99% confidence level, Aharonian et al. 2005)
to above 400 GeV (using
from the original publication) yields
I(>400 GeV
cm-2 s-1.
4 Coordinated XMM-Newton/RXTE observations
PKS 2005-489 was observed three times with XMM-Newton between 2004 and 2005
(Obs-Id 0205920401, 0304080301 and 0304080401).
The satellite pointings were scheduled such that simultaneous HESS
observations were possible. The first XMM-Newton pointing
was performed on October 4, 2004, and yielded a net exposure of 11 ks.
Unfortunately poor weather did not allow HESS observations on this night.
Good quality HESS data were taken on the following nights.
Two more pre-planned XMM-Newton observations were performed on the nights of
September 26 and 28, 2005,
with exposures of 21 and 25 ks, respectively. The goal
of scheduling such temporally close
XMM-Newton pointings was to sample both the shortest variability time scales
and spectral variations occurring on the typical time scale of HBL
(one to a few days; see e.g. Tanihata et al. 2001). It should also be noted that the
simultaneous observations had to occur near the end of the HESS observing
season for PKS 2005-489, because of the narrow overlap between the
HESS and XMM-Newton visibility windows caused by the
constraints on the orientation of the solar panels of the satellite.
During the XMM-Newton observations, all three EPIC CCD cameras were used along
with the thin filter. The PN detector (Strüder et al. 2001) was operated in timing mode,
i.e. when the data from a predefined area of the CCD chip are collapsed into a one-dimensional
row to be read out at high speed.
This mode allows the sampling of the shortest possible flux and spectral variations
without potential pile-up
problems during bright states.
The two MOS (Turner et al. 2001) instruments were used in different configurations.
MOS1 took data in timing mode in 2004 and in small-window mode during both 2005
observations. The MOS2 camera was operated in small-window mode in 2004
and in large-window mode in 2005.
Simultaneous observations were also performed with the Optical Monitor (OM) onboard XMM-Newton,
using all photometric filters sequentially. For these data the central window was read
in fast mode to enable temporal studies within the data. The exposure for each filter
varied between 1800 and 4400 s, ultimately constrained by the overall duration of the pointing.
In August 2005, target-of-opportunity observations of PKS 2005-489 with RXTE were triggered based on apparent enhanced VHE activity in preliminary analysis of uncalibrated HESS data. The satellite pointings were scheduled such that further HESS data could be taken simultaneously. Unfortunately, only a limited subset of the HESS data in this epoch passes standard quality selection criteria because of poor weather conditions in Namibia.
4.1 XMM-Newton data analysis
The EPIC data were processed and analyzed with SAS v7.1.0, using the calibration files as of July 2008 and standard screening criteria![[*]](/icons/foot_motif.png)




RAWX


RAWX

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Figure 4:
Time-average VHE energy spectrum observed from PKS 2005-489.
The dashed line represents the best |
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![]() |
Figure 5:
Photon spectra observed by HESS
from PKS 2005-489 during each year of data taking.
Only the statistical errors are shown.
Each line represents the best |
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The spectral analysis was performed with XSPEC v11.3.2ag, testing different binning schemes with at least 50 counts in each new bin. For the 2004 observation, the spectrum from the MOS1 instrument is not included since the timing mode data are significantly noisier than, but consistent with, the other detectors. The MOS and PN spectra are fit together, with a free constant to allow for different MOS/PN normalizations. The PN flux is adopted as the reference value, but the MOS fluxes are typically within a few percent of the PN fluxes. The X-ray spectra are each fit with single and broken power-law models, with interstellar absorption modelled by TBabs (Wilms et al. 2000). This model is used with the solar abundances of Wilms et al. (2000) and the cross-sections by Verner et al. (1996) (for a discussion, see Baumgartner & Mushotzky 2006). The results of the fits are shown in Fig. 7 and Table 3 (details in Sect. 5.2).
The data from the OM were processed with the
tasks omichain for the photometry and omfchain for the timing analysis.
Since no variations are found in the light curves,
the whole exposure is used to derive the flux measurements.
The photometry was processed interactively with omsource for every filter,
to assure the use of point-source analysis procedures.
A standard aperture of 6 pixels on the binned images
(12 for the unbinned images of the 2004 October data set) is used for all filters.
This corresponds to an aperture of 6
for the optical filters (V, B, U).
The counts for the UV filters are extrapolated by the software, using the UV point-spread
function, to an aperture of 17
5.
These are the two apertures for which the OM count-rate-to-flux conversion is calibrated.
The OM images are affected by stray light from a bright star
in the field of view, which increases the background near one side
of the source region. Therefore background events are taken from both an
annulus around the source and from two different circular regions
at the same distance of the source from the main stray-light reflex.
Tests performed with different background sizes and locations
show that the photometry does not change by more than 1%,
well below the systematic uncertainties of the flux conversion
(estimated at
10%, see XMM-SOC-CAL-TN-0019).
The source fluxes were obtained from the count rates using the
OM conversion factors for white dwarfs,
adding the 10% systematic error in quadrature.
The fluxes were de-reddened for Galactic absorption
using the extinction curve by Cardelli et al. (1989) with the updates by O'Donnell (1994),
and assuming RV[=
AV/E(B-V)]=3.1. This is the average value for the Galactic diffuse ISM.
For the line of sight of PKS 2005-489, a value of
is used
(from NED; Schlegel et al. 1998), corresponding to
AV(5500)=0.182.
The conversion factors, extinction ratios and the resulting source fluxes
are reported for each of the filters in Table 4.
PKS 2005-489 is hosted by a giant elliptical galaxy of
total R magnitude 14.5 and half-light radius 5.7
(from HST snapshot observations, Scarpa et al. 2000).
For the SEDs (Figs. 9 and 10),
the OM fluxes were corrected for the contribution of the host galaxy.
The wavelength-dependent correction was determined using a template
SED for elliptical galaxies (Silva et al. 1998),
rescaled to the host-galaxy flux in the R band, and accounts
for the given apertures. The OM fluxes are always dominated
by the non-thermal emission, and the small contribution of
the host-galaxy is only noticeable in the V and B filters.
![]() |
Figure 6: VHE spectrum measured by HESS from PKS 2005-489 in 2004 compared to the previously published version (Aharonian et al. 2005). Only the statistical errors are shown. Correcting the 2004 data for decreases in the optical efficiency of the HESS array results in a three-times higher integral flux above 400 GeV. |
Open with DEXTER |
![]() |
Figure 7: Best-fit data and folded model, plus residuals, of the XMM-Newton PN ( upper) and MOS2/1 ( lower) spectra of PKS 2005-489, in October 4, 2004 ( left panel) and September 26, 2005 ( right panel). The spectra are fitted with a concave broken power-law model plus Galactic absorption. |
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Table 3: Fits to the X-ray spectraa in the different epochs.
4.2 RXTE data analysis
The RossiXTE/PCA (Jahoda et al. 1996) performed 15 snapshot observations of PKS 2005-489
between August 20 and September 19, 2005, yielding a total exposure of 23.4 ks.
The observations were mostly done with PCU0 and PCU2.
The PCA STANDARD2 data were reduced and analyzed
with the standard FTOOLS routines in HEASOFT V6.3.2,
using the filtering criteria recommended by the RXTE Guest Observer Facility.
The new SAA history file and parameters for the background calculation provided
in September 2007 were used,
which correct a bug in the calculation of the faint background
model.
Only the top-layer events were processed and only the PCU2 data were considered,
to ensure a more accurate spectral measurement.
The average net count rate in the 3-15 keV band is
cts s-1 pcu-1.
The RXTE spectra were extracted and fitted separately for each pointing,
and summed together to obtain the average spectrum.
Each spectrum is well-fit by a single power-law model.
5 XMM-Newton/RXTE results
5.1 Flux variability
The 2004 and both 2005 EPIC light curves
contain no evidence of flux variability within each exposure, and
have average net count rates in the MOS2 camera of
,
,
and
counts s-1, respectively.
The probability of constant emission is
,
0.92 and 0.99
for the three epochs, respectively. Similarly, no
flux variability is observed in the OM data for any of the observed filters.
No significant flux variations (
)
are detected
during the RXTE observations (see Fig. 8) as well. However,
night-by-night flux variations
of a few tens of percent
cannot be excluded, due to low statistics in the RXTE measurements.
Although no variability is found within any of the XMM-Newton exposures
or within the complete RXTE sample, strong flux variations
are found on longer time scales. The 2-10 keV flux varies by a factor of 16 between 2004 and 2005. In the UV and optical bands the flux increases by
40% and
20%, respectively, between 2004 and 2005.
A smaller (factor of 2.5), but significant, change in the 2-10 keV flux
also occurs during the 17 days between the final RXTE and first XMM-Newton pointing in September 2005.
5.2 X-ray spectrum
Table 4: Optical Monitor parameters and source fluxes for the three XMM-Newton observations.
![]() |
Figure 8:
Simultaneous X-ray and VHE data taken by RXTE, XMM-Newton and HESS
in August-October 2005. Upper panel: the X-ray photon index. The XMM-Newton values were obtained from fits above 2 keV, to be consistent with the RXTE energy range.
Lower panel: the observed X-ray and VHE flux.
The X-ray data (filled circles) are energy fluxes integrated in the 2-10 keV band.
The HESS data (open squares) are integral photon fluxes above 400 GeV.
The vertical scales on the left and right are for the X-ray and HESS data, respectively.
All errors are shown at the 1 |
Open with DEXTER |





In all three data sets, the EPIC spectra are better fit by a broken power-law model
with respect to a single power law, with high significance (F-test >99.9%).
Remarkably, the parameters reveal an inverted broken power-law spectrum,
where the slope in the soft X-ray band is steeper than in the hard X-ray band
(
).
![]() |
Figure 9:
SED of PKS 2005-489 in the frequency range
covered by the XMM-Newton and RXTE observations (filled circles). Archival X-ray spectra
are shown along with those
measured by XMM-Newton on Oct. 4, 2004 (red points), RXTE in Aug.-Sept. 2005 (green points)
and XMM-Newton on Sept. 26, 2005 (blue points). For the latter,
two fits are plotted, showing the effects of using two different
estimates of the Galactic column density
(
|
Open with DEXTER |
In the 2004 dataset, the X-ray spectrum is characterized by the lowest flux
ever recorded from this source (see Fig. 9)
and a very steep slope of
up to
3 keV.
At higher energies, there is evidence of a moderate hardening trend.
The trend is consistent between the MOS and PN spectra,
but is only significant in the PN data thanks to better
statistics.
This hardening can also be reproduced by the sum of two
power-law models, the second having a harder index.
The slope of this second power-law function is not well constrained.
However, assuming a fixed value of
,
as could be expected for the inverse Compton emission of low-energy electrons,
a good fit (
for 283 NDF) is obtained with
a flux normalization at 1 keV for the second power-law function
at
% of the first power law.
In this case, the two power-law functions would cross at about 20 keV.
Using the high
has minimal impact on this result,
steepening the low-energy slope by only +0.1.
The difference between the EPIC spectra
of September 26 and 28, 2005 is negligible. However, the spectrum
from either of the 2005 observations is significantly
harder and brighter than in 2004 (Fig. 9).
Either 2005 spectrum is represented well by a pure power-law model with
from
1 to 10 keV.
At lower energies, each has a slightly steeper slope.
Although modest (
), the spectral break
is significant (F-test >99.99%)
and is not caused by the extension of the MOS data at low energies,
or by the cross-calibration between the detectors. The broken power-law model
is also statistically required for the PN spectrum alone (F-test >99.99%),
as well as for the MOS spectra fitted separately.
It should be noted however that the precise location of the break
is affected by systematic more than statistical errors.
In Table 3, the break energy for the fit with
the overall minimum
is reported, but there are other similar
local minima between 0.7 and 2.2 keV,
whose hierarchy in
can change according to the calibration
of the effective area near the instrumental Si and Au edges (at 1.7 and 2.1 keV).
The true break therefore should be considered more realistically inside the range 0.7-2.2 keV.
This does not affect the values of the soft and hard slopes in a relevant way
(they do not change by more than 0.04 and 0.02, respectively).
The amount of the spectral break, however, depends on the adopted value of
the Galactic column density. Using the DL value instead
of the LAB survey value, the soft X-ray slope becomes steeper (
),
making the break more pronounced (
).
The two cases are shown in Fig. 9,
where the spectra are plotted together with the other data sets
and archival data. Allowing instead the column density to be lower than all
available estimates, a single power-law spectrum is obtained with
cm-2 (
for 859 NDF), though a broken power-law
model still seems to provide a better representation (F-test >95%).
![]() |
Figure 10: SED of PKS 2005-489, in different epochs and states. The XMM-Newton and RXTE data are plotted as in Fig. 9, together with historical data (shown in grey; from Massaro et al. 2008; Tagliaferri et al. 2001). The recent Fermi-LAT spectrum (0.2<E<10 GeV) from Aug.-Oct. 2008 is also plotted, for reference (open bow-tie; Abdo et al. 2009). For clarity, the HESS spectra (E>300 GeV) are plotted as bow-ties, and have been corrected for EBL absorption with the model by Franceschini et al. (2008, see Table 5). Main panel: the HESS time-averaged spectrum from 2004 (red filled bow-tie) and 2005 (blue hatched bow-tie). Right mini-panel: the HESS spectrum in 2005, divided in two subsets: before the XMM pointings (periods 07-08/2005 in Fig. 3; green filled bow-tie) and during the XMM epoch (period 09/2005 in Fig. 3; blue hatched bow-tie). Same colors associate optical, X-ray and VHE datasets corresponding to same epochs. |
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The 2005 RXTE observations yield a very steep (
)
spectrum from 3-15 keV.
This spectrum is almost parallel to the 2004 XMM-Newton spectrum but with
6 times higher flux. The RXTE photon index and flux are
in good agreement with the values measured by Swift in April 2005 (Massaro et al. 2008).
This suggests that the hardening of the X-ray spectrum of PKS 2005-489
occurred in September 2005, between the RXTE and XMM-Newton pointings.
6 Discussion
The spectral energy distribution (SED) of PKS 2005-489,
assembled with the data taken in 2004 and 2005, is shown in Fig. 10.
The VHE -ray spectra are corrected for the energy-dependent attenuation
caused by interactions with the optical-IR photons of the diffuse extragalactic
background light (EBL).
The correction was performed adopting the EBL model of Franceschini et al. (2008).
This model takes into account the most recent results on galaxy emission and evolution,
and agrees with both the upper limits on the EBL derived from blazar spectra
(Aharonian et al. 2007,2006c) and the lower limits given by galaxy counts,
at both optical and infrared wavelengths (Madau & Pozzetti 2000; Fazio et al. 2004).
The fits to the absorption-corrected spectra are reported in Table 5,
and plotted in Fig. 10.
Table 5: Best power-law fits to the HESS spectra corrected for EBL absorption following Franceschini et al. (2008).
The simultaneous observations do not show evidence
of strong changes in the location
of the SED peaks over the years, with respect to the historical values.
The hard optical-UV spectrum from the OM photometry
and the steep X-ray spectrum above 0.1 keV
locate the synchrotron peak between the two bands, at approximately
Hz.
At VHE, the steep spectra constrain the peak of the
-ray emission
to energies <0.2 TeV.
The low flux in the VHE band does not enable a determination of the spectral properties
in exactly the same (short) observing windows of the X-ray observations.
However, given the lack of significant VHE variability in 2004 or 2005,
the monthly or yearly average spectra can be considered a reasonable approximation
to the VHE spectra during the epoch of the X-ray observations.
In 2004, PKS 2005-489 was in a very low state in both the X-ray and VHE bands.
The spectral slopes are also similar (
)
for the X-ray and VHE emission.
The X-ray flux around few keV is more than two orders of magnitude less than
the BeppoSAX flux during the active state of 1998. The slight spectral hardening
towards higher X-ray energies is consistent with the onset of the concavity
expected from the emergence of the inverse-Compton emission in the SED.
This emission, produced by low-energy electrons, is usually not visible
in the X-ray spectra of HBL because of the dominance of
the synchrotron emission from high-energy electrons. In this data set,
its presence becomes detectable because of the steepness and very low flux
of the synchrotron (X-ray) spectrum. A similar behavior has recently been observed
in the HBL PKS 2155-304 during a low state (Zhang 2008; Foschini et al. 2008).
This transition between steep synchrotron and flat
inverse-Compton spectra
is instead fully visible
inside the observed X-ray passband for intermediate BL Lac objects
(e.g. ON 231, see Tagliaferri et al. 2000).
In 2005, the X-ray spectrum changes radically.
Swift and RXTE measurements made in April and August-September 2005, respectively,
show a large increase in flux but little change in the spectral slope
compared to 2004. Shortly after the RXTE observations, a
much harder spectrum (
)
and a brighter state than in the previous epochs is found during XMM-Newton observations.
Most interestingly, the data require a spectral break with a steeper spectrum
towards lower energies. This break is more apparent
when a higher galactic
is used
(e.g. the DL versus LAB survey values; see Fig. 9),
but even when considering a single power-law shape with an ad-hoc
,
the
extrapolation of the X-ray spectrum down to lower frequencies does
not match the UV-optical data, which are still rising with frequency.
A hump with a steeper soft X-ray slope is needed to smoothly
connect with the optical-UV spectrum, and this feature
indicates the presence of two particle populations, since it
cannot be easily reproduced by synchrotron radiation with a single electron distribution.
At VHE, the comparison between the 2004 and 2005 spectra indicates
a hardening trend as well,
though with marginal significance (1 sigma; see Table 5)
due to the large errors on the HESS measurement.
However, if this trend is real, and if the VHE spectrum is indeed produced by the same
particles as are responsible for the X-ray emission
(as suggested by the similar spectral slopes
in 2004-2005 and by the location of the synchrotron
peak
),
a similar behavior is expected. Namely, the VHE spectral hardening
should have occurred after the RXTE observations, during the epoch of the XMM-Newton pointings,
while the 2005 July-August VHE spectrum should be similar to the 2004 spectrum (modulo normalization).
To test this hypothesis, the 2005 HESS data set was divided into two subsets.
The first consists of all VHE data before the XMM-Newton pointings
(2005 July-August dark periods),
and the second contains all 2005 HESS data taken during or after the pointings
(2005 September dark period).
Table 5 shows the results of power-law fits to the spectra from these subsets,
after correction for the absorption of VHE photons on the EBL.
The HESS spectrum in the XMM epoch is indeed harder ()
than the pre-XMM spectrum, and the statistical significance of the spectral
change increases (
2
)
with respect to the comparison between yearly spectra.
In addition, the pre-XMM spectrum has the same slope as the 2004 HESS spectrum.
The HESS data therefore are fully consistent with a behavior in the VHE band
that mirrors the spectral variations seen in the X-ray band,
as expected if both emissions are produced by the same electrons.
Interestingly, the harder VHE state continues in 2006.
Because of the limited
exposure in 2007, the VHE statistics are too poor
to draw any meaningful conclusions.
The most remarkable feature of the VHE emission, however, is the lack of
flux variability. Comparing the average states in 2004 and 2005,
the integrated energy flux (e.g. for reference in the decade 0.3-3 TeV)
does not vary by more than 40% between 2004 and 2005,
and is consistent with a constant flux. In contrast, the X-ray flux
increased by a factor of 16 in the 2-10 keV band (see Table 3).
The optical-UV emission, which is close to the synchrotron peak
and which typically provides the target photons for the IC scattering
in the VHE band, shows an increase as well (
40%).
In a homogeneous SSC scenario, if the VHE emission is produced
by the same electrons emitting in the X-ray band
(as indicated in this case by the spectral behavior),
a fresh injection of electrons should make the VHE flux increase
at least linearly with the X-ray flux.
The reason is that the energy density of all possible seed photons
for the IC scattering has increased or is constant between the two epochs.
For the VHE flux to remain constant, a corresponding decrease in the seed photons
energy density is required.
The X-ray and VHE data taken together, therefore, suggest that a new jet component is emerging in the SED of PKS 2005-489, which is physically separated from the main emitting region. The emission at the synchrotron peak of this new component should be lower, remaining hidden below the observed SED, while its harder spectrum at high energies emerges in the hard-X band. The electrons of this new component would not see the energy density of the observed synchrotron peak, but the lower energy density of their self-produced synchrotron peak instead.
7 Conclusions
PKS 2005-489 is detected at VHE in each of the four
years it was being observed by HESS (2004-2007). The 2005-2007 data clearly confirm the VHE discovery reported by
HESS (Aharonian et al. 2005) and quadruple the statistics
of the initial spectrum measurement. Re-analysis of the previously
published 2004 HESS data, using the improved calibration of the detector's energy
scale, results in a 3 times higher flux and a similar photon index.
The measured VHE -ray flux is low (
3% Crab)
and only shows weak variations on time scales ranging from days to years.
The observed time-averaged (2004-2007) VHE spectrum is
soft, with a photon index
.
Although evidence of VHE spectral variations is marginal by itself,
the VHE spectrum seems to track the X-ray slope variations when
multi-wavelength coverage is available.
Observations performed with XMM-Newton and RXTE in 2004 and 2005
reveal remarkable changes in the X-ray spectrum,
but without shifting the location of the synchrotron
peak with respect to historical observations.
Interpreting these measurements along with the HESS data suggests the emergence of a new jet component in the SED that is characterized by a harder electron spectrum. This component must be separate: its particles cannot interact with the synchrotron photons of the observed SED peak, otherwise higher VHE fluxes than observed would be implied.
PKS 2005-489 is found overall in a very low state,
in both the X-ray and VHE bands, during the observations
presented in this article. PKS 2005-489 has historically
demonstrated a 100
dynamical range in the X-ray band.
Thus, dramatically higher VHE fluxes (
)
can be expected in the future, unless such an increase
in the X-ray flux is counter-balanced by a strong (>10
)
and simultaneous increase in the blazar's magnetic field.
Further monitoring of this object is highly encouraged,
as it is one of the few HBL easily detected at VHE during
low states and has the potential for extreme brightness
and variability.
The results presented here confirm the strong diagnostic potential
of coordinated optical - X-ray - VHE observations. Future studies
can be significantly improved by incorporating data from the recently launched
Fermi -ray satellite. The Fermi data will
provide information on the lower energy side of the inverse-Compton peak,
enabling contemporaneous measurements of both sides of each blazar hump.
The support of the Namibian authorities and of the University of Namibia in facilitating the construction and operation of HESS is gratefully acknowledged, as is the support by the German Ministry for Education and Research (BMBF), the Max Planck Society, the French Ministry for Research, the CNRS-IN2P3 and the Astroparticle Interdisciplinary Programme of the CNRS, the UK Science and Technology Facilities Council (STFC), the IPNP of the Charles University, the Polish Ministry of Science and Higher Education, the South African Department of Science and Technology and National Research Foundation, and by the University of Namibia. We appreciate the excellent work of the technical support staff in Berlin, Durham, Hamburg, Heidelberg, Palaiseau, Paris, Saclay, and in Namibia in the construction and operation of the equipment. The article is based on observations obtained with XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. The authors thank the RXTE team for support during the ToO trigger. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory.
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Footnotes
- ...
- Supported by CAPES Foundation, Ministry of Education of Brazil.
- ...
- Now at Harvard-Smithsonian Center for Astrophysics, Cambridge, USA.
- ...
- Now at W.W. Hansen Experimental Physics Laboratory & Kavli Institute for Particle Astrophysics and Cosmology, Stanford University, Stanford, USA.
- ... France
- UMR 7164 (CNRS, Université Paris VII, CEA, Observatoire de Paris).
- ... index
- The power-law spectrum is described as
.
- ... consistent
- The
difference between
the two positions is
.
- ... published
- The flux was reported
above a threshold of 200 GeV. Extrapolating the earlier result,
using
as reported, yields I(>400 GeV) =
cm-2 s-1.
- ... close
- The blazar was observed in 2 consecutive orbits.
- ... pile-up
- That is, when more than one X-ray photon arrives in one camera pixel or adjacent pixels before the CCD is read out.
- ... criteria
- see SAS User Guide and CAL-TN-0018.
- ... dwarfs
- http://xmm2.esac.esa.int/sas/7.1.0/watchout/
- ...
model
- Details at http://www.universe.nasa.gov/xrays/programs/rxte/pca/doc/bkg/bkg-2007-saa/
- ...
peak
- If the VHE electrons were emitting by synchrotron at energies well below the X-ray range, they would correspond to the synchrotron peak, around 1016 Hz. Since the SSC scattering at the peak would be in the Thomson regime for this frequency, the gamma-ray peak should then be inside the VHE band, which is not observed.
All Tables
Table 1: Results from long-term HESS observations of PKS 2005-489.
Table 2:
Best
fits of a power-law to the various
spectra of PKS 2005-489 measured by HESS.
Table 3: Fits to the X-ray spectraa in the different epochs.
Table 4: Optical Monitor parameters and source fluxes for the three XMM-Newton observations.
Table 5: Best power-law fits to the HESS spectra corrected for EBL absorption following Franceschini et al. (2008).
All Figures
![]() |
Figure 1:
Distribution of |
Open with DEXTER | |
In the text |
![]() |
Figure 2: Integral flux, I(>400 GeV), measured by HESS from PKS 2005-489 during each dark period of observations (i.e. moonless night time within a month). Only the statistical errors are shown. The horizontal line represents the average flux for all the HESS observations. The four horizontal line segments represent the average annual flux observed during the corresponding years (see Table 1). |
Open with DEXTER | |
In the text |
![]() |
Figure 3: Integral flux, I(>400 GeV), measured by HESS from PKS 2005-489 during each night of observations. The individual plots represent each of the four years (2004-2007) of data taking. Only the statistical errors are shown. The horizontal lines represent the average annual flux observed during the respective year. The vertical lines at MJD 53282, 53640, and 53642 represent the nights of XMM-Newton observations. The epoch of the RXTE observations is between the vertical lines at MJD 53609 and 53623. |
Open with DEXTER | |
In the text |
![]() |
Figure 4:
Time-average VHE energy spectrum observed from PKS 2005-489.
The dashed line represents the best |
Open with DEXTER | |
In the text |
![]() |
Figure 5:
Photon spectra observed by HESS
from PKS 2005-489 during each year of data taking.
Only the statistical errors are shown.
Each line represents the best |
Open with DEXTER | |
In the text |
![]() |
Figure 6: VHE spectrum measured by HESS from PKS 2005-489 in 2004 compared to the previously published version (Aharonian et al. 2005). Only the statistical errors are shown. Correcting the 2004 data for decreases in the optical efficiency of the HESS array results in a three-times higher integral flux above 400 GeV. |
Open with DEXTER | |
In the text |
![]() |
Figure 7: Best-fit data and folded model, plus residuals, of the XMM-Newton PN ( upper) and MOS2/1 ( lower) spectra of PKS 2005-489, in October 4, 2004 ( left panel) and September 26, 2005 ( right panel). The spectra are fitted with a concave broken power-law model plus Galactic absorption. |
Open with DEXTER | |
In the text |
![]() |
Figure 8:
Simultaneous X-ray and VHE data taken by RXTE, XMM-Newton and HESS
in August-October 2005. Upper panel: the X-ray photon index. The XMM-Newton values were obtained from fits above 2 keV, to be consistent with the RXTE energy range.
Lower panel: the observed X-ray and VHE flux.
The X-ray data (filled circles) are energy fluxes integrated in the 2-10 keV band.
The HESS data (open squares) are integral photon fluxes above 400 GeV.
The vertical scales on the left and right are for the X-ray and HESS data, respectively.
All errors are shown at the 1 |
Open with DEXTER | |
In the text |
![]() |
Figure 9:
SED of PKS 2005-489 in the frequency range
covered by the XMM-Newton and RXTE observations (filled circles). Archival X-ray spectra
are shown along with those
measured by XMM-Newton on Oct. 4, 2004 (red points), RXTE in Aug.-Sept. 2005 (green points)
and XMM-Newton on Sept. 26, 2005 (blue points). For the latter,
two fits are plotted, showing the effects of using two different
estimates of the Galactic column density
(
|
Open with DEXTER | |
In the text |
![]() |
Figure 10: SED of PKS 2005-489, in different epochs and states. The XMM-Newton and RXTE data are plotted as in Fig. 9, together with historical data (shown in grey; from Massaro et al. 2008; Tagliaferri et al. 2001). The recent Fermi-LAT spectrum (0.2<E<10 GeV) from Aug.-Oct. 2008 is also plotted, for reference (open bow-tie; Abdo et al. 2009). For clarity, the HESS spectra (E>300 GeV) are plotted as bow-ties, and have been corrected for EBL absorption with the model by Franceschini et al. (2008, see Table 5). Main panel: the HESS time-averaged spectrum from 2004 (red filled bow-tie) and 2005 (blue hatched bow-tie). Right mini-panel: the HESS spectrum in 2005, divided in two subsets: before the XMM pointings (periods 07-08/2005 in Fig. 3; green filled bow-tie) and during the XMM epoch (period 09/2005 in Fig. 3; blue hatched bow-tie). Same colors associate optical, X-ray and VHE datasets corresponding to same epochs. |
Open with DEXTER | |
In the text |
Copyright ESO 2010
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