4.3 Data analysis and results
4.3.3 Spectroscopy: looking for the reflection imprints
Method
In order to study the spectral features typical of an XRN, we have considered the PN camera spectra of the regions highlighted in Fig.4.5 (see also Table 4.2). For this analysis we selected only PN spectra because of the much higher effective area in the hard X- ray band (&6 keV). Building the PN spectra for the Fe-Kα bright filaments, we used
the same GTI filtered event files as employed to build the fluence map in Section 4.2. We stacked all the spectra (of both the source and the background) and the response files using the ftools MATHPHA, ADDARF and ADDRMF, without caring about the 6.4-keV line variability. The background region was centred on R.A.=17:45:42.289, DEC.=-29:07:53.39 (circle 1 arcmin in radius), at the same Galactic latitude as the filaments in order to avoid any systematics due to the different intensity of the Galactic Ridge emission at different Galactic latitudes. The spectral model used in the fitting comprises two thermal plasmas (APEC; Smith et al. 2001) with a metallicity fixed at 2 times solar, two gaussian emission lines (GAUSS) which account for the Fe-Kα and Kβ lines, and a non-thermal continuum
(POWER LAW). All of these components are subject to photoelectric absorption in the line-of-sight column density (WABS; Morrison & McCammon 1983, common to all the spectral components). The ratio between the Fe Kβ and Fe-Kα line fluxes has been fixed
at 0.11 (Koyama et al., 2009). Moreover, to quantify the presence of a putative Fe-K edge, we multiplied the hard power-law continuum by an EDGE component at a fixed energy of 7.1 keV. We note that, although we have applied the background subtraction technique in this study, all the spectra still exhibit a residual contribution from the hot
APEC component (kT∼6.5 keV). We have therefore retained this spectral component in the fitted model, a different approach to the one we adopted in our study of the MCs in the
98 line reverberation in Galactic Centre molecular clouds
Arches cluster region (Capelli et al., 2011b). Although the temperature of this hot plasma is approximately the same all over the inner CMZ, its intensity can significantly change; the best evidence of the substructures of this spectral component is the relative contribution of the Fe XXV Kα line at 6.7 keV in the spectra shown in Fig.4.6; the MCs A, B1, DS1
and DS2 show the significant presence of the 6.7 keV line, while in the spectra of the other regions the contribution from this fluorescent line is lower. To avoid high residuals in the spectral domain of the Fe fluorescent lines (6-7 keV), we therefore accounted for this spectral component in the fitting model of all the MCs. In the spectral fitting procedure we used again the Cash statistic because of the low signal-to-noise ratio in the spectral channels.
The specific goal of this analysis is to quantify the depth (τ) of the K edge at 7.1 keV and also to establish the equivalent width (EW) of the 6.4-keV line with respect to the underlying ionising continuum (POWER LAW). Our investigation of the spectral characteristics of the various regions and sub-regions is more detailed than that recently reported by Ponti et al. (2010).
Results
Sunyaev & Churazov (1998) have shown that in the XRN model, the 6.4-keV line emission has to be accompanied by strong absorption above 7.1 keV, the minimum energy that a photon must have in order to be able to produce a K-shell vacancy in neutral Fe via the photoelectric effect. Moreover, because the primary source of ionisation is not seen by the observer, the EW of the line is expected to be very high, i.e. of the order of 1 keV. On the other hand, the depth of the Fe absorption edge at 7.1 keV is strongly related to the intrinsic NH of the MC, with cloud column densities in excess of 1023cm−2 imprinting
edge features with optical depth of about 0.1-0.2 at 7.1 keV on the emergent reflected (i.e., electron-scattered) continuum. In Fig.4.6 and Table 4.5, we present the results of our search for the imprints of reflection on the time averaged spectra of the selected GC molecular filaments. We measured a very high value of the EW of the 6.4-keV line in all the MCs we studied; the largest values have been measured in the MCs B2 and F, where we also see a variable the Fe-Kα line flux (respectively a decreasing and increasing flux).
We first note that the NH values for the clouds reported Ponti et al. (2010), as inferred
from radio CS measurements, are always lower than 1023cm−2. Specifically, we measured the strongest absorption edge in the spectrum of region A (called MC1 in Ponti et al., 2010), where the Fe optical depth has been measured to be 0.6±0.1. This is rather peculiar since this cloud has one among the lowest values of NH (4×1022cm−2 as inferred from the
CS measurements) among the MCs considered. On the other hand, MCs with higher reported column densities show no Fe-K edge at 7.1 keV. For example, the regions B2 and D (subregions of the larger complex referred to by Ponti et al. (2010) as the bridge) have (9×1022cm−2, see Table 2 in Ponti et al., 2010), but we could only measure an upper limit
to the Fe absorption edge.
In Section 4.3.5 we present a prescription to calculate the NH of the MCs from the
4.3 Data analysis and results 99
Table 4.5: Results from the analysis of the PN time averaged spectra of the Fe-Kα bright
filaments. The spectral parameters reported in the table are the column density ( in units of 1022cm−2), the temperatures of the warm and the hot plasma components (in keV), the
normalizations of the warm and hot plasmas (in units of 10−17 and 10−18Rn
enHdV/4πD2,
respectively), the slope of the powerlaw with its normalization (in units of 10−5 pho- tons keV−1 cm−2 s−1 at 1 keV), the peak energy and theσ of the Fe-K
α line (in keV), the
6.4-keV line flux (in units of 10−5 photons cm−2 s−1), the EW of the 6.4-keV line (in keV),
the optical depth at the Fe-K edge and the C-stat value for the best fit model.
A B1 B2 C D NH 6.7+5−0..34 7.7 +0.6 −0.4 7.6±0.4 7.9±0.5 6.4±0.4 kTwarm 0.7±0.1 0.9±0.1 0.9±0.1 0.8±0.1 1.0±0.1 norm 3.7+6−0..19 4.4+1−0..97 3.2−+00..96 4.5+1−1..20 2.10.4 −0.5 kThot 6.7±1.2 6.4+1−1..56 5.1 +1.8 −1.9 6.6 +2.0 −2.2 6.5 fix norm 2.2+0−0..45 2.0−+00..24 1.1−+00..43 1.4±0.3 0.6±0.3 Γ 0.62+0−0..2405 1.21+0−0..1209 1.1±0.1 1.3+0−0..12 1.0+0−0..23 norm 6.8+3−1..81 14.2+5−2..51 7.4+2−1..74 12.0+5−2..39 4.0+2−1..07 E6.4 6.420+0−0..007004 6.416 +0.007 −0.006 6.415 +0.003 −0.004 6.421 +0.006 −0.013 6.417 +0.011 −0.007 σ 0.03±0.01 0.05±0.01 0.039+0−0..005008 0.042+0−0..015011 0.043+0−0..012020 F6.4 1.86+0−0..87 1.42 +0.07 −0.06 1.61±0.06 0.98 +0.06 −0.05 0.87±0.05 EW 0.86+0−0..2132 0.94+0−0..2329 1.64+0−0..4640 0.98+0−0..3631 1.36+1−0..1151 τ 0.6±0.1 0.2±0.1 ≤0.07 0.23+0−0..1519 ≤0.4 C-stat 1842.62 1742.87 1568.51 1731.63 1691.50 E F DS1 DS2 NH 6.5+0−0..43 7.8±0.2 8.1 +0.9 −1.0 9.2 +1.6 −1.1 kTwarm 1.0±0.1 1.08−+00..0203 0.6±0.1 0.4±0.1 norm 4.1+0−0..87 22.5+1−1..89 8.6+5−4..47 42.2+3−8..42 kThot 6.5 fix – 7.1+2−1..00 6.6 +1.2 −0.9 norm 0.9±0.5 – 3.1−+00..76 5.4+1−0..06 Γ 1.0±0.2 1.1±0.2 0.8±0.3 1.0+0−0..32 norm 9.0+3−2..99 39.7+18−12..70 6.4+3−2..96 18.6+13−6.1.7 E 6.4 6.417±0.008 6.411±0.005 6.414+0−0..010009 6.412±0.007 σ 0.030+0−0..018014 0.029+0−0..012009 0.041+0−0..019014 0.053+0−0..009012 F6.4 1.62±0.08 7.5±0.3 1.17+0−0..0807 2.1±0.1 EW 1.23+0−0..6641 1.54+0−0..5280 0.84+0−0..8134 0.77+0−0..4335 τ 0.22+0−0..2120 0.34−+00..2524 0.4±0.2 0.53+0−0..1418 C-stat 1751.11 1843.80 1653.91 1761.42
100 line reverberation in Galactic Centre molecular clouds 10−3 0.01 normalized counts s −1 keV −1
The A molecular cloud
5 0 0.5 1 1.5 ratio Energy (keV) 10−3 0.01 normalized counts s −1 keV −1
The B1 molecular cloud
5 0 0.5 1 1.5 ratio Energy (keV) 10−4 10−3 0.01 normalized counts s −1 keV −1
The B2 molecular cloud
5 0 0.5 1 1.5 ratio Energy (keV) 10−3 0.01 5×10−4 2×10−3 5×10−3 0.02 normalized counts s −1 keV −1
The C molecular cloud
5 −1 0 1 2 χ Energy (keV) 10−4 10−3 0.01 normalized counts s −1 keV −1
The D molecular cloud
5 0.5 1 1.5 ratio Energy (keV) 10−3 0.01 normalized counts s −1 keV −1
The E molecular cloud
5 0 1 2 ratio Energy (keV) 10−3 0.01 0.1 normalized counts s −1 keV −1
The F molecular cloud
5 0 1 2 ratio Energy (keV) 10−3 0.01 normalized counts s −1 keV −1
The DS1 molecular cloud
5 0.5 1 1.5 ratio Energy (keV) 0.01 2×10−3 5×10−3 0.02 0.05 normalized counts s −1 keV −1
The DS2 molecular cloud
5 −1 0 1 ratio Energy (keV)
Figure 4.6: Time-averaged PN spectra of the set of MCs in the present study. Each panel shows the measured data along with the best fitting spectral model and (in the lower section) the ratio of these two quantities . The residuals which can be seen in the spectra of the regions A, B1 and DS2 are due to the inhomogeneous pattern of the Cu fluorescent line at 8.05 keV in the XMM PN detector. This are systematics due to the background subtraction, which however do not influence the measurements of the other spectral parameters, namely the optical depth at the Fe-K absorption edge and the slope of the non-thermal power law.
generally give us higher column densities for the MCs than the ones assumed by Ponti et al. (2010).
A very important result is the identification of a hard non-thermal component associ- ated with the 6.4-keV line, which in the XRN model represents the fraction of incident reflected (i.e. Thomson scattered) by the molecular material. In all the observed spectra we were able to discern underlying rather hard continuum modelled as a power-law with a photon index in the range Γ=0.6-1.3. The different determinations of Γ are not consistent
4.3 Data analysis and results 101
with a single value within the 90% uncertainties but this may be result of the impact of varying effective NH values for the clouds (see Section 4.3.5).
The value Γ=1.3 determined for cloud C is the highest amongst the set of measure- ments. Interestingly in March 2004 region C also contained the X-ray transient XMMU J174554.4-285456, whose best fit model during the outburst activity was found to be an absorbed power law (NH=1.41×1023cm−2 and Γ=1.86, Porquet et al., 2005b). It is cer-
tainly plausible that reflection by Cloud C of the light from the transient accounts for the upward biasing of the slope of the hard continuum in the time-averaged measurement.
We have also investigated the characteristics of the Fe-Kα line, in terms of its energy
and width. For this purpose we built a time-averaged spectrum for the whole region, stacking all the spectra across all the observations. In this stacked spectrum, the intrinsic width of the line is 36±3 eV, while its peak is at 6414±2 eV, slightly higher than the nominal value of 6403 eV for neutral iron1. The results from the individual spectra (see Table 4.5) show the same trend, with the measured Fe-Kαline energy in all cases exceeding
the theoretical value. Although the EPIC calibration is known to be better than 10 eV, the presence of soft protons in the FOV might slightly change the energy and the width of the lines (Ponti et al., 2010). To check whether these are instrumental effects or real features, we measured the same quantities for the Cu instrumental line at 8.04 keV. The centroid of the Cu Kα line is at 8047+4−2 eV, while its width is 34+6−4 eV. The line peak is in good
agreement with its nominal value of 8048 eV. Although a relative shifting between the two lines is present (the Cu Kα peak perfectly matches the theoretical expectation, whereas
the Fe-Kα is noticeably higher), taking into account the systematics and the statistical
uncertainties, we conclude that no statistically significant line blueshift and broadening have been measured and that our results are consistent with emission from neutral, or close-to-neutral, Fe atoms embedded in the material of the cold molecular clouds. We also note that Suzaku measured the Fe-Kα line energy in the molecular filaments to be 6409±1
eV with a line width of 33+2−4 eV, the latter being marginally higher than the expected systematics (≈30 eV) (Koyama et al., 2007a).