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Precise Chandra Localization of the Soft Gamma-ray Repeater SGR 1806-20

We present observations of the Soft Gamma-ray Repeater SGR 1806-20 taken with the Chandra X-ray Observatory. We identify the X-ray counterpart of SGR 1806-20 based on detection of 7.5-s pulsations. Using three unrelated X-ray sources (and USNO stars) as position references, we are able to determine that the SGR is at alpha_{2000}=18h08m39.32s and delta_{2000}=-20o24'39.5", with rms uncertainty of 0.3"in each coordinate. We find that SGR 1806-20 is located within the 1-sigma error region determined by Interplanetary Network data and is 14+/-0.5 arcsec distant from the non-thermal core of SNR G10.0-0.3, excluding SGR 1806-20 as the origin of the core. We see evidence for a significant deviation of the spin-down of SGR 1806-20 from its long-term trend, demonstrating erratic spin-down behavior in this source similar to that seen in other SGRs. Finally, we show that there is a broad X-ray halo surrounding SGR 1806-20 out to radii ~1' due to scattering in the interstellar medium.

uncertainty of 0. ′′ 3 in each coordinate. We find that SGR 1806−20 is located within the 1-σ error region determined by Interplanetary Network data and is 14±0.5 arcsec distant from the non-thermal core of SNR G10.0−0.3, excluding SGR 1806−20 as the origin of the core. We see evidence for a significant deviation of the spin-down of SGR 1806−20 from its long-term trend, demonstrating erratic spin-down behavior in this source similar to that seen in other SGRs. Finally, we show that there is a broad X-ray halo surrounding SGR 1806−20 out to radii ∼ 1 ′ due to scattering in the interstellar medium.

Introduction
The soft γ-ray repeaters (SGRs; see Hurley 2000 for a recent observational review) are enigmatic sources that were discovered through their repeated bursts of hard X-rays/soft γrays. With the improvement in imaging afforded by the current generation of X-ray satellites SGRs were found to emit softer quiescent X-ray emission as well. They are generally thought to be young (< 10 4 yr) neutron stars with extremely strong magnetic fields, or magnetars (e.g. Duncan & Thompson 1992;Thompson & Duncan 1993). This belief has been motivated by their associations with young supernova remnants (SNRs; Evans et al. 1980;Kulkarni & Frail 1993;Vasisht et al. 1994) and sites of massive star formation (Fuchs et al. 1999;Vrba et al. 2000), the energetics of their bursts (Paczyński 1992;Thompson & Duncan 1995), and the detection of X-ray pulsations with long (5-10 s) periods and large (∼ 10 −11 s s −1 ) spin-down rates (e.g. Kouveliotou et al. 1999).
The proposed identification of a SNR, the plerion G10.0−0.3, with SGR 1806−20 by Kulkarni & Frail (1993) played a key role in establishing the connection of SGRs with SNRs and thence young stars. This identification renewed confidence in the association of N49 with the 5 March 1979 event (SGR 0526−66) and motivated SNR associations for two of the remaining SGRs, SGR 1900+14 ) and SGR 1627−41 (Hurley et al. 1999b). The association of SGR 1806−20 with G10.0−0.3 was strengthened by the discoveries of a non-thermal core in G10.0−0.3 Vasisht, Frail, & Kulkarni 1995) that changed on month to year time-scales and of a nearby luminous blue variable star (LBV; Kulkarni et al. 1995;van Kerkwijk et al. 1995). It was proposed that all of these sources were related, with the LBV as either a current or past binary companion to SGR 1806−20  and the variable radio source the result of episodic mass loss from SGR 1806−20 (Frail, Vasisht, & Kulkarni 1997).
However, a recent moderate-precision IPN position (Hurley et al. 1999a) placed SGR 1806−20 outside the non-thermal core suggested by Kulkarni et al. (1994) to be the seat of the SGR. Furthermore, the association of SGR 1900+14 with SNR G42.8+0.6 was questioned (Lorimer & Xilouris 2000). In summary, the entire issue of whether SGRs are associated with SNRs is now open to debate (see review by Gaensler et al. 2001).
The purpose of this paper is to provide an independent localization of the quiescent X-ray counterpart of SGR 1806−20. SNR/SGR associations are useful for establishing independent ages for SGRs; the non-thermal core to G10.0−0.3 provides key energetics constraints; and the LBV star hints at a binary origin and evolutionary scenario for SGR 1806−20. Determining whether or not these associations are true is therefore key to assessing the nature of SGR 1806−20.

Observations & Analysis
We observed SGR 1806−20 with the Chandra X-ray Observatory on two occasions, beginning on 2000 July 24.7 UT and 2000 August 15.7 UT, with durations of 4.9 ks and 31 ks respectively (see Table 1). Both observations were taken with the aim point on the backside-illuminated ACIS S-3 detector. The first observation was taken in the standard full-frame CCD mode with 3.24-s time resolution, while the second observation was acquired in a 1/4-subarray mode that provides increased time resolution of 0.81 s, but with a reduced field-of-view (1/4 of the full-frame area). All the data suffer from the effects of photon pileup: ≈ 50% of the data from the first observation and < ∼ 10% of the data from the second are corrupted. Pile-up occurs when "two or more photons are detected as a single event" 1 . This results in 3 main effects: (1) the detected event has an energy that is the sum of all of the incident photon energies; (2) the count rate is diminished; and (3) the grades of the incoming photons will be altered, so that some may be cataloged as "cosmic rays" or other undesireable events and will therefore be rejected. The consequences of these effects on an observation of a bright source, in addition to modification of the spectrum, are that the spatial profile is altered, as the diminishment of the count-rate is worst for the pixels at the center of the psf, and that some regions with especially high rates of incoming photons may appear devoid of all events, as the total amplitude of the detected events exceeds the onboard energy threshold and is rejected (similar to grade migration).
Because of pile-up, especially effects (1) and (2) above, we cannot perform spectroscopy with great accuracy and we defer such analysis until tools for correcting pile-up become publicly available. As noted, pile-up will also affect the spatial profile of the source (it depresses the center of the psf relative to the wings) and the timing (two photons arriving together will be counted as one) but the analysis presented here should not be affected significantly, as we are not interested in the detailed shape of the psf within 1 ′′ or in the exact lightcurve -overall, the position and period will be preserved.
We examined the two images and find a single bright source near the expected location of the SGR. This source is consistent with an unresolved point source (Gaussian σ = 0. ′′ 33; Marshall et al. 2001). In order to determine if this source is indeed the SGR, we searched for the expected 7-s pulsations (Kouveliotou et al. 1998). We first barycentered the data using the axBary software. The 5-ks observation had too few photons to detect anything, and it was taken with 3.24 s time resolution. The 31-ks data had sufficient photons and was taken in the 1/4 subarray mode, which gives 0.81 s sampling. We therefore added a random number ∼ U(0, 0.81) to the time-of-arrival values to eliminate any effects of sampling. We then performed a Z 2 1 test (de Jager, Raubenheimer, & Swanepoel 1989) on the 31-ks data, and find a very significant periodicity at P ≈ 7.5 s. To refine this measurement, we performed a phase connection of the Chandra data in the manner of Fox et al. (2001). We connected four segments of ≈ 8000-s, each of which had been binned to 1-s resolution (appropriate for the 0.8-s sampling). We find, referenced to MJD 51772.0 (TDB), a phase of 0.12(5) cycles, a period P = 7.4925(2) s, and a period derivative |Ṗ | < 4 × 10 −8 s s −1 at 90% confidence. The pulsations have a sinusoidal profile and a rms pulsed fraction of ≈ 7.4% -a folded pulse profile is shown in Figure 1.
Although the count-rate and spectrum from our observations are uncertain due to pileup, we can roughly estimate the source flux through use of previously published spectra for SGR 1806−20 (Mereghetti et al. 2000). With the count-rate from the 31-ks observation (the least corrupted), we see 0.5-10 keV fluxes of ≈ 7 × 10 −12 erg s −1 cm −2 (absorbed) and ≈ 2 × 10 −11 erg s −1 cm −2 (unabsorbed). These values are entirely compatible with the 2-10 keV fluxes found by Mereghetti et al. (2000). The observed count-rates are constant throughout the observations. Given the 0.81-s sampling and the presence of pile-up, it is not surprising that we do not see any significant bursts of the type typically seen during observations of SGRs (e.g. Fox et al. 2001). We note that it is likely any bursts during this observation were probably not recorded due to pile-up and to on-board/post-processing rejection of anomalous events.

Localization
To accurately localize SGR 1806−20, we measured its position in both Chandra datasets. The measured positions from the two observations are consistent to 0. ′′ 04, suggesting that the stochastic position errors are minimal. There may however be systematic errors on the order of 1 ′′ due to overall uncertainties in the CXO aspect solutions.
To correct for such errors we searched the Chandra images for background sources to use as position references. We found seven weak sources in the 5-ks image. None had a match in the SIMBAD database, but two sources (detected at 2.5-σ confidence) were coincident with stars. In the 31-ks image we found 11 sources, two of which had matches in the USNO-A2.0 catalog: one detected at 4-σ confidence that was also in the 5-ks image, the other detected at 2.5-σ confidence. We summarize the X-ray detected USNO stars in Table 2. There are ≈ 4.3 USNO stars arcmin −2 in this region, giving chance coincidence rates of ∼ 0.1% between an individual X-ray source and a USNO star, supporting our identifications. We are pursuing photometric and spectroscopic observations of these stars that should allow us to classify them and verify these identifications (e.g. van den Berg & Verbunt 2001).
This position agrees very well (within the 1-σ error ellipse) with the position determined from IPN measurements (Hurley et al. 1999a), and is 14 ± 0.5 ′′ from the position of the nonthermal nebula from Frail et al. (1997); see Figure 2.

Extended Emission
X-rays are scattered by dust grains in the interstellar medium and one expects to see halos of size 0.1-100 arcmin towards objects with hydrogen column density of 10 21 cm −2 or greater (i.e., Overbeck 1965;Predehl & Schmitt 1995). This issue was first discussed by Overbeck (1965) and the Einstein Observatory found the first such halos (Rolf 1983;Catura 1983). The most comprehensive work to date has been carried out by Predehl and associates using the ROSAT Observatory (Predehl & Schmitt 1995).
The recent interest of scattering halos around SGRs has been motivated by the discovery of such a halo around SGR 1900+14 (Kouveliotou et al. 2001). As with SGR 1900+14, one expects an X-ray halo around SGR 1806−20 from scattering off dust. We therefore examined the radial profile of SGR 1806−20 from the 31-ks observation for better signal-to-noise. In Figure 3, we show the radial profile of SGR 1806−20 from this observation along with the profile from a MARX 3 simulation (using the spectrum for SGR 1806−20 from Mereghetti et al. 2000) scaled to the same normalization for the inner portion (cf. Kouveliotou et al. 2001); such simulations should be accurate out to radii of a few arcminutes 4 . The background was measured from the ACIS S-3 chip far (> 3 ′ ) from the SGR 1806−20, and includes both instrumental and unresolved-source contributions. We corrected the background estimate to r < ∼ 1 ′ with a MARX simulation of the ACIS S-3 response to account for vignetting etc. We find the variation in radial response to be small for r < ∼ 1 ′ , and therefore show a constant background in Figure 3. One sees very good agreement between the data and the model at small radii, showing that pile-up has not substantially corrupted the radial profile of this observation (the model did not incorporate pile-up, as MARX does not correctly simulate the back-illuminated CCDs). But there is a clear deviation from the model psf at large radii ( > ∼ 3 ′′ ) that we believe to be indicative of a dust halo contributing ∼ 0.5% to the total X-ray flux. To verify this interpretation we examined the radial profile as a function of the phase of the SGR 1806−20. The profile for the portion of the period when SGR 1806−20 is "ON" has of course more counts at small radii than the profile when SGR 1806−20 is "OFF." However, the outer portion (r > ∼ 3 ′′ ) that we identify as the halo is identical between "ON" and "OFF", indicating that the pulsations of SGR 1806−20 are smeared out in the halo. This suggests that the extended emission is not an instrumental property, as scattering wings or something similar would scale with the immediate flux of the source instead of averaging over time as observed here. We note, though, that while the data do show a slight softening towards higher radii (as expected from the E −2 dependence of the scattering cross-section; e.g. Predehl et al. 2000), we cannot determine in detail how the halo profile changes with energy. This could be due to the relatively hard spectrum of SGR 1806−20, so that the cross section does not vary much over the energy range for which there are significant counts (2-6 keV). It could also be an effect of poor statistics, where a longer observation would show an energy dependent halo. Regardless, we must caution the reader as to this interpretation of the extended emission.

Discussion
The X-ray position of SGR 1806−20 is now without doubt incompatible with the positions of both the non-thermal core of G10.0−0.3 and the LBV star, although those two positions still agree to 1 ′′ . We therefore follow Hurley et al. (1999a) and Gaensler et al. (2001) in suggesting that the LBV, not SGR 1806−20, powers the core of G10.0−0.3. As outlined in Hurley et al. (1999a) the LBV can easily explain the energetics and changing morphology of the non-thermal core. SGR 1806−20 could still be associated with G10.0−0.3 as a whole, although this is unlikely given the proximities of the core and the LBV to the center of G10.0−0.3. However, we still have a remarkable coincidence: within a circle ≈ 10 ′′ in radius (0.7 pc at 14.5 kpc distance; Corbel et al. 1997), we have an SGR (SGR 1806−20), an LBV, and a massive star cluster (Fuchs et al. 1999). All three of these objects are exceedingly rare. In addition, there is a SNR 5 at the same position. While it is hard to draw quantitative conclusions from this coincidence since none of the objects involved have well-defined populations and this region of the sky has been studied in great depth, it is still noteworthy. Individual distances to these sources are not known, but it is plausible that they are all 14.5 kpc away (Corbel et al. 1997;Fuchs et al. 1999). We believe that even if SGR 1806−20 is not directly associated with the LBV or G10.0−0.3, it was likely born in the same cluster, which requires that the progenitor of SGR 1806−20 was quite massive to have had a supernova before the LBV or the stars in the cluster (as noted by Hurley et al. 1999a). Perhaps this is a case of an extremely high-mass star ( > ∼ 50M ⊙ ) forming a neutron star due to large mass loss over its lifetime (Fryer & Kalogera 2001)? Alternatively, the sources could all have originated in the same molecular cloud, but the supernova from SGR 1806−20 could have triggered the star formation that lead to the LBV and the cluster.
The angular scale where the scattering halo departs from the psf is smaller than that seen for ROSAT data, and we do not measure the halo out to such large radii, so the relations presented for ROSAT data by Predehl & Schmitt (1995) are not directly applicable. However, we can correct the relation between fractional halo intensity and hydrogen column density from Predehl & Schmitt (1995) to the appropriate energy and angular scale range for our data, and we find an expected halo intensity of ≈ 0.3% based on the N H = 6×10 22 cm −2 , as determined by Mereghetti et al. (2000). This is very similar to the measured intensity, supporting the conclusion that we see a halo and not an instrumental artifact, and implying that the size distribution of the scattering grains at small angular scales is similar to that seen for the larger ROSAT halos.
The halo we see is very similar (in angular scale and fractional intensity) to that seen for SGR 1900+14 (Kouveliotou et al. 2001), suggesting that they lie behind similar dust columns (or possibly that both "halos" are in fact instrumental effects). As with SGR 1900+14, the pulsations of SGR 1806−20 are too smeared out in the halo to use them to determine a geometric distance (cf. Trümper & Schönfelder 1973;Predehl et al. 2000). However, if observations were made of SGR 1806−20 soon (∼ hours to days) after a significant change in the level of its emission (e.g. following a giant flare) such that the change was visible in the halo, a distance determination could be made. BeppoSAX is probably the only current satellite that has the necessary slew capabilities coupled with spatial resolution and softenergy response to perform these observations. The measured period of SGR 1806−20 is not consistent with either its long-term spindown Mereghetti et al. 2000) or the newer spin-down measured from phase-connected RXTE observations , as seen in Figure 4. While SGR 1806−20 exhibits substantial timing noise , the long-term trend was generally stable for a number of years. However, Woods et al. (2000) found that at a single epoch the instantaneous spin-down rate for SGR 1806−20 was higher by a factor of 1.5 that the global trend, and to account for the new Chandra data, the spin-down rate must have increased even more. Between the last RXTE measurement from Woods et al. (2000) and our measurement, the average spin-down rate is 2.3(1) × 10 −10 s s −1 , which is a factor of 2 higher than the largestṖ previously measured for this source. This change is similar to the large change in both spin period and spin-down rate seen for SGR 1900+14 following its activation during 1998 August (Marsden, Rothschild, & Lingenfelter 1999;Woods et al. 1999) and for the anomalous X-ray pulsar 1E 1048.1−5937 (Kaspi et al. 2000). SGR 1806−20 must have either increased its long-term spin-down substantially or suffered from a momentary large change in spin-down. Such sudden changes in spin-down are capable of being produced within the magnetar framework , either due to particle outflows or due to re-alignment of the neutron star crust, but they are typically assumed to be episodic and triggered by bursting or flaring activity Thompson et al. 2000). We note that Gögüş et al. (2000) did observe a period of moderate burst activity for SGR 1806−20 several days before the 31-ks Chandra observation, but this was too close to the Chandra observation to have affected the spin-period significantly. Whether sudden or gradual, the change in rotation for SGR 1806−20 had to start before the bursts observed by Gögüş et al. (2000) (unless the spin-down is orders of magnitude larger than previously seen), so these changes were not caused by burst activity. This suggests that the more gradual change is the correct mechanism in this case, perhaps due to continuing plastic deformation of the crust (e.g. Thompson et al. 2000). In addition, spin-down cannot be used to determine the ages or magnetic fields of any SGR in the manner typical for radio pulsars (e.g . Manchester & Taylor 1977), as noted by Thompson et al. (2000). We must instead rely upon arguments relating to the super-Eddington bursts (Thompson & Duncan 1995), global energetics, and similar phenomena as indicators of the magnetar-like fields of these objects.

Conclusions
We have confirmed the position of Hurley et al. (1999a) for SGR 1806−20, and determined that it is near neither the non-thermal core of G10.0−0.3 (Frail et al. 1997) nor the luminous blue variable star (LBV; van , although those two sources may be associated. Through comparison of the radial profile of SGR 1806−20 to a model psf, we see evidence for a broad X-ray scattering halo. Given its violent nature, SGR 1806−20 may produce changes in its X-ray emission such that observations of this halo could lead to a geometric distance determination, thereby fixing the luminosity scale of SGRs and locating SGR 1806−20 in three dimensions relative to the massive star cluster (Fuchs et al. 1999), the LBV, and other nearby objects.
Timing analysis of the Chandra data show that SGR 1806−20 is not continuing its longterm trend of spin-down that it has followed for the past 4 years. This is similar to changes in spin-down observed for SGR 1900+14 following the 1998 August outbursts, although only minor bursts were observed for SGR 1806−20 around the time of these observations and these bursts are unlikely to have been responsible for the deviation from normal spin-down.
We would like to thank R. Rutledge and J. Carpenter for valuable conversations. D.L.K. is supported by the Fannie and John Hertz Foundation, S.R.K. by NSF and NASA, and E.V.G. and G.V. by the NASA LTSA grant NAG5-22250. We have made extensive use of the SIMBAD database, and would like to express our appreciation of the astronomers who maintain this database.   (Frail et al. 1997) with position of radio core at square marked "CORE". Superimposed are 1-, 2-, and 3-σ annuli around the best-fit IPN position (asterisk labeled "IPN") from Hurley et al. (1999a), the corrected Chandra position (circle labeled "CXO"), and the position of the LBV star (triangle marked "LBV"; Kulkarni et al. 1995). Note that the position of the core indicated in Hurley et al. (1999a) is slightly incorrect (K. Hurley 2001, personal communication). Adapted from Hurley et al. (1999a).