XMM-Newton Users Handbook

3.5.5 OM sensitivity and detection limits

Both, Fig. 106 and Table 20, provide the expected count rate of the OM, in counts per second, for $m_V = 20$ mag stars of various types, with the different filters introduced above. These numbers were obtained under the assumption of a perfect detector, i.e., free of deadtime and coincidence loss, and without time sensitivity degradation, and with an aperture radius of $6''$ ($17''.5$) for optical (UV) filters.

Table 20: OM count rates [$10^{-4}$ counts s$^{-1}$ ] as function of spectral type for stars with $m_V~=~20$ mag under the assumption of a detector with zero deadtime, no coincidence loss and no time sensitivity degradation.

Filter	B0	A0	F0	G0	G2	K0	M0	WD
V	1555	1510	1473	1473	1495	1452	1392	1412
B	6190	5272	3805	3010	2577	2289	1337	4454
U	8533	2216	1486	1119	966	444	178	4667
UVW1	5647	958	366	218	182	29.1	17.4	3093
UVM2	2276	316	39.3	8.69	5.54	0.436	0.032	1174
UVW2	950	126	11.1	2.06	1.19	0.112	0.069	561

Figure 106: OM count rates vs. filter selection for stars of different spectral type with $m_v = 20$ mag.

The numbers listed in Table 20 can be used to calculate the perfect detector count rates of stars of given magnitude by the formula

$$C_{mag} = C_{20} \times 10^{0.4\times(20-b_{mag})} \quad ,$$ (3)

where $C_{20}$ is the count rate of an $m_V = 20$ mag star of the same spectral type as the target, and $b_{mag}$ and $C_{mag}$ are the $m_V$ and the count rate of the target of interest, respectively.

However, OM deadtime and coincidence losses must be taken into account, for instance using SAS. For the 20th magnitude stars in Fig. 106 and Table 20, OM coincidence losses are negligible. Losses become significant for a point source at a count rate of about 10 counts s$^{-1}$ (about 10% coincidence) for a full frame exposure.

The correction is approximated by the following formula:

$$C_{mag} = \frac{\ln (1 - C_{det}*T)}{T_{ft}-T} \quad$$ (4)

where $C_{mag}$ is the count rate of incident photons from equation (3), $C_{det}$ is the actual measured count rate, $T$ is the CCD frametime in units of seconds and $T_{ft}$ the frametransfer time in the same units (0.1740 ms). The frametime is a function of the OM science window configuration, as described in § 3.5.2.

An optimum aperture radius of $6''$ was found to be the radius at which the loss correction formula is self consistent. At this radius the accuracy of the coincidence loss correction is at a level of 5% between different frametimes.

The validity range of the coincidence loss correction can be obtained from eq. 4 as a function of the frametime, whose maximum value (11.04 ms) occurs for full frame exposures, high or low resolution, where the whole detector is used. In windowed exposures the frametime varies from about 5 to 10 ms. For this reason the maximum measured count rates that can be reliably corrected are in the range of a few hundreds counts per second. Higher rates (1000 cts/s) may not damage the detector, but are scientifically useless.

We give in Table 21 coincidence loss corrected rates for limit high measured rates. The table gives the maximum measurable rate for each given value of frametime (1 count per frametime) and the corresponding corrected values for the maximum rate minus 0.01 and 0.001 c/s. For these high rates we obtain correction factors of 10, and very small differences in count rate produce variations in the corrected rates of more than 20%.

Table 21: OM coincidence loss corrections for high count rates.

Frametime	Maximum	Corrected rate	Corrected rate
	measured rate	max-0.01	max-0.001
(ms)	(c/s)	(c/s)	(c/s)
5	200	2052	2531
7	143	1402	1739
9	111	1056	1317
11	91	842	1054

The coincidence loss correction implemented in SAS (from version 14) sets a validity limit at a rate of 0.97 counts per frametime (source+background). Within this limit, the correction error is less than 2%. This corresponds to about 700 and 300 source+background corrected counts for frametimes of 5 and 11 ms, respectively. For values higher than 0.97, the source counts are not corrected by SAS, and therefore, the conversion to magnitudes and fluxes are not applied. The produced source list contains only the measured count rate and position for these sources.

In Fig. 107, a comparison between formula (4) and the inflight instrument performance is shown.

Figure 107: Theoretical (red line) and empirical (blue line) coincidence loss correction curves superposed to flight data

In Table 22, an estimate of the limiting magnitudes that can be detected with OM in an integration time of 1000 s at the 5-sigma confidence level is provided. They have been obtained from simulations, taking into account the whole aperture ($6''$ radius for optical filters and $17''.5$ for UV) on which the OM photometric calibration is based. If we consider that detections can be obtained in the very few pixels in the core of the PSF, since the total background is then lower, the limit would be about one magnitude fainter.

Table 22: Limiting magnitudes$^1$ for a 5-$\sigma$ detection in 1000 s

Filter	Spectral type
	B0	A0	G0	K0	WD
V	19.8	19.8	19.7	19.7	19.7
B	21.0	20.8	20.2	19.9	20.6
U	21.8	20.4	19.6	18.6	21.2
UVW1	21.1	19.2	17.6	15.4	20.5

Notes to Table 22:
1) Numbers obtained adopting as background a zodiacal light level of 100 S10 units, where S10 is the brightness of one m$_V$=10 mag solar type star per square degree. All magnitudes are in V filter.

Table 23 presents the limiting magnitudes derived from the photometric analysis of one of the OM calibration fields (HD5980, also observed from the ground). This can be considered as an extreme case because of the crowdedness of the field which increases the observed background by overlap of the PSF wings. This is why the obtained limits are brighter than the simulations of Table 22.

Table 23: Observed limiting magnitudes$^1$ for a 5-sigma detection (scaled to 1000s)

Filter	Spectral type range
	B0/A0	A1/G3	G4/M8
V	19.5	19.4	19.4
B	20.6	19.8	19.3
U	20.2	19.4	18.6

Notes to Table 23:
1) Based on OM observations of HD5980 field, obtained with V, B and U filters. Magnitudes are standard V.

A detailed study of the background sources in OM can be found in:
http://www.cosmos.esa.int/documents/332006/623312/MSSLbckgd_v6.pdf
http://www.cosmos.esa.int/documents/332006/623312/OMbackg05.pdf.
The expected levels of different external background radiation processes in the optical/UV are tabulated in Table 24. The background count rate in the OM is dominated by the zodiacal light in the optical. In the far UV the intrinsic detector background becomes important. Images are regularly taken with the blocked filter and no LED illumination to measure the detector dark counts. The mean OM dark count rate is 4.0 $\times 10^{-4}$ counts s$^{-1}$ pixel$^{-1}$. The variation across the detector is $\pm$ 9% , with a mainly radial dependence, being highest in an annulus of about 8' radius and lowest at the centre. Variations as a function of time are broadly cyclic, on a timescale of 11 years (likely associated to the solar cycle), with a semi-amplitude of 23%. If a very bright star is in the field of view, the dark rate can be up to $\simeq$60% higher than normal, despite the use of the Blocked filter.

Table 24: Levels of different OM background contributors

Background source	Occurrence	Count rate range
Diffuse Galactic$^1$	all directions	2.14$\times10^{-1}$-7.527$\times10^{-6}$
Zodiacal$^{1,2}$	longitude $90^\circ\pm20^\circ$	1.69$\times10^{-1}$-5.611$\times10^{-6}$
Average dark count rate$^3$	all directions	$4.0 \times 10^{-4}$

Notes to Table 24:
1) In units of [counts s$^{-1}$ arcsec$^{-2}$].
2) In differential ecliptic coordinates, i.e., the angle between the Sun and the pointing position of XMM-Newton. Maximum is for l,b = 70,0 in the white filter, minimum is for l,b = 110,70 in the UVW2 filter.
3) In units of [counts s$^{-1}$ pixel$^{-1}$]

Artifacts can appear in the XMM-Newton OM images due to light being scattered within the detector. These have two causes: internal reflection of light within the detector window and reflection of the off-axis starlight and background light from part of the detector housing.

The first of these causes a faint, out of focus ghost image of a bright star, displaced in the radial direction away from the primary image (see Fig. 108).

Figure 108: Out-of-focus ghost image (“smoke ring”) of a bright field star, during a 3C273 observation. Clearly visible is the strong fixed pattern noise around the bright source

The second effect is due to light reflecting off a chamfer in the detector window housing. Bright stars that happen to fall in a narrow annulus $12'.1$ to $13'$ off axis shine on the reflective ring and form extended loops of emission radiating from the centre of the detector (see Fig. 109).

Figure 109: Straylight ellipses caused by reflection of a star outside the field of view, taken from the PKS 0312 offset 6 field (V filter in use). The average background rate is 15 counts pixel$^{-1}$ per exposure; in the bright straylight loop is 30 counts pixel$^{-1}$. The background is also enhanced in the central region due to reflection of the diffuse sky light from outside the field. In the centre it rises to $\simeq$3 times the average background rate.

Similarly, there is an enhanced “ring” of emission near the centre of the detector due to diffuse background light falling on the ring. These features are less prominent when using UV filters, due to the lower reflectivity of UV light.

Figure 110 shows the low sensitivity patch produced by an accidental observation of Jupiter in July 2017. This figure shows also the "ring" mentioned before. The low sensitivity patch is oval with a size around $210 \times 120$ pix$^2$, equivalent to 1.4 square arcmin approximately. The whole field is 17 arcmin in diameter.

Figure 110: Image obtained with the V filter showing the depletion patch caused by Jupiter in July 2017. The RGS and EPIC-pn boresights are marked in red.