How can I convert from OM count rates to fluxes?
Any XMM-Newton users,
who need to convert OM count rates in the optical and UV filters into
fluxes (expressed, e.g., in erg/cm/cm/s/A), can do it by
following any of the methods we outline below.
OM count rates can be obtained with SAS. The chain omichain allows any user to process OM image data. Among other
products, it produces source list files where the count rates for all
objects detected in an image are given. These rates, as from SAS 6.5,
are
corrected from all known instrumental effects, mainly coincidence
losses, dead time and time sensitivity degradation. The count rates can
also be obtained from the PPS pipeline products delivered by the SSC at
Leicester.
In case a user wants to measure the count rates by him/herself, we
provide a recipe at the end of this document.
Method 1 (flux conversion based in white dwarf standard stars
observations)
- multiply the count rate obtained with SAS (files
*SWSRLI*, corrected count rate in CORR_RATE column) by the
corresponding average conversion
factors listed in the following Table.
| Filter |
Effective
wavelength(nm) |
Conversion factor
(erg/cm/cm/A/cnt) |
| V |
543 |
2.49E-16
|
| B |
450 |
1.29E-16
|
| U |
344 |
1.94E-16
|
| UVW1 |
291 |
4.76E-16
|
| UVM2 |
231 |
2.20E-15
|
| UVW2 |
212 |
5.71E-15
|
These conversion factors have been
derived from observations of white dwarf standard stars. They give an
average transformation which can be used in all cases where no
information is available on the spectral characteristics of the
source. The error of the fluxes obtained by using these factors is less
than 10%.
In case of stellar objects for which the spectral type is known, the
user may prefer to use Method 2.
Method 2 (spectral type dependent flux conversion)
- choose the most appropriate spectral type (in case of white
dwarf, refer to Method 1)
- convert from count rate, obtained with SAS (files
*SWSRLI*, corrected count rate in CORR_RATE column), to flux, selecting
the most
appropriate factor for the chosen spectral type and filter
from the following Tables:
| Filter |
A0V |
B0V |
F0V |
G0V |
K0V |
M0V |
Vega |
| V |
2.50E-16 |
2.48E-16 |
2.52E-16 |
2.54E-16 |
2.56E-16 |
2.65E-16 |
2.50E-16 |
| B |
1.36E-16 |
1.16E-16 |
1.41E-16 |
1.53E-16 |
1.60E-16 |
1.81E-16 |
1.34E-16 |
| U |
1.71E-16 |
1.94E-16 |
1.80E-16 |
1.83E-16 |
1.88E-16 |
2.01E-16 |
1.70E-16 |
| UVW1 |
4.96E-16 |
4.72E-16 |
4.96E-16 |
4.51E-16 |
3.88E-16 |
1.09E-16 |
4.86E-16 |
| UVM2 |
2.20E-15 |
2.14E-15 |
2.10E-15 |
1.84E-15 |
1.66E-15 |
n.a.
|
2.19E-15 |
| UVW2 |
6.06E-15 |
5.56E-15 |
7.15E-15 |
6.05E-15 |
5.76E-16 |
n.a.
|
5.88E-15 |
These numbers have been obtained by folding the
spectral library of Pickles with the in-flight response curves
of the OM. For Vega we have used the alpha_lyr_stis_002 calibrated
spectrum provided by the HST Calspec database.
If we compare these factors with the ones derived from white dwarfs
(Method 1) we see that they are very similar except for the UV filters
in cool stars. These is mostly due to the fact that the library spectra
for these stars have very low or zero flux at ultraviolet wavelengths.
Similar results are obtained with a different library: the
Bruzual-Persson-Gunn-Stryker Spectra:
| Filter |
A0V |
B0Ib |
F0IV |
G0V |
K0V |
M0V |
| V |
2.48E-16 |
2.50E-16 |
2.50E-16 |
2.55E-16 |
2.56E-16 |
2.61E-16 |
| B |
1.29E-16 |
1.17E-16 |
1.38E-16 |
1.44E-16 |
1.55E-16 |
1.80E-16 |
| U |
1.66E-16 |
1.97E-16 |
1.77E-16 |
1.88E-16 |
1.85E-16 |
1.94E-16 |
| UVW1 |
4.79E-16 |
4.76E-16 |
4.84E-16 |
5.02E-16 |
5.15E-16 |
3.14E-16 |
| UVM2 |
2.15E-15 |
2.17E-15 |
2.18E-15 |
2.27E-15 |
2.02E-15 |
1.42E-15 |
| UVW2 |
5.56E-15 |
5.25E-15 |
6.14E-15 |
6.50E-15 |
6.34E-15 |
2.46E-15 |
As in the case of Pickles library, the UV part of the spectra for
cool stars is not reliable.
A calibration observing campaign is currently being carried out with
OM in order to obtain the most reliable conversion factors for each
stellar type. The users are invited to check periodically the content
of this
page, to ensure that the most updated calibrations are always
being employed.
Method 3 (AB magnitude system)
The OM Team has implemented the AB magnitude system for OM.
Data processed with SAS in the SSC general re-processing, or using a
version higher than 6.5 will contain these magnitudes and also the
corresponding fluxes in the combined source list file.
The AB system is defined for OM as:
mAB(filter) = Zero_point(filter)
- 2.5 log10(Count_Rate)
where the Zero points for each filter come from the original definition
of AB magnitude as:
Zero_point(filter) = -48.60 - 2.5
log10(1/n_phot)
with n_phot being the
count rate measured in a filter for a constant incoming flux of one
erg/sec/cm2/hz. The calculated zero points for the OM AB system are
then:
| Filter |
Zero-point |
| V |
17.923 |
| B |
19.081 |
| U |
19.189 |
| UVW1 |
18.566 |
| UVM2 |
17.412 |
| UVW2 |
16.572 |
The inverse of n_phot provides
a count rate to flux conversion factor (in frequency space) for each
filter:
| Filter |
Conversion factor
(erg/cm/cm/hz/cnt) |
| V |
2.46E-27
|
| B |
8.47E-28
|
| U |
7.66E-28
|
| UVW1 |
1.36E-27
|
| UVM2 |
3.94E-27
|
| UVW2 |
8.54E-27
|
These conversion factors are obtained by folding a constant flux in
frequency space of one erg/sec/cm2/hz with the in-flight response
curves
of the OM.
Method 4 (General method. Flux based on Vega flux scale)
This describes the whole process. Steps 1 to 7 are currently performed
by SAS. The magnitudes of the detected objects are written in the
source lists for each exposure (*SWSRLI* file) and in the final
combined list (*OBSMLI* file)
- determine the source+background counts within an aperture of 12
pixels radius (unbinned) for U, B and V filters. For UVW1, UVM2 and
UVW2, the counts measured within the 12 pixels aperture are
extrapolated to 35 by using the PSF definition. Users doing this "by
hand" should take an aperture of 35 pixels. In case of binned images
(2x2) aperture sizes shoudl be reduced by half.
- determine the background counts within an annulus of radii
14 and 25 pixels (unbinned) for U, B and V filters (8 and 13 in binned
images). For UVW1, UVM2 and UVW2 the inner radius should be at least 37
- apply the coincidence loss and deadtime corrections
to the source+background counts as
specified in the
XMM-Newton User Handbook
- subtract the background counts
- correct for time sensitivity degradation as described in these Evergreen tips and
tricks
- calculate the instrumental magnitude, using the formula:
mag = -2.5 log10 (count rate) + zero-point
where the zero-points are reported in the following table:
| Filter |
Zero-point |
| V |
17.963 |
| B |
19.266 |
| U |
18.259 |
| UVW1 |
17.204 |
| UVM2 |
15.772 |
| UVW2 |
14.867 |
This is the way SAS uses to derive the
magnitude of any detected star in an OM image. The use of SAS is
strongly recommended.
- Vega is used as a reference to determine the magnitude to flux
conversion. The flux of Vega should correspond to a magnitude
0.03 in OM B,V filters, and 0.025 in OM U and UV filters. The Vega
fluxes in the OM filters are listed in the following Table:
| Filter |
Effective
wavelength (nm) |
Flux (erg/cm2/s/A) |
| V |
543 |
3.70E-09 |
| B |
450 |
6.40E-09 |
| U |
344 |
3.20E-09 |
| UVW1 |
291 |
3.68E-09 |
| UVM2 |
231 |
4.33E-09 |
| UVW2 |
212 |
5.03E-09 |
The quoted fluxes for Vega have been
obtained by interpolation in its calibrated spectrum provided by the
HST Calspec database.
- apply the following formula for a given filter
mag (Vega) - mag (source) = -2.5 log
[Flux(Vega)/Flux(source)]
Pages maintained by SAS
librarian.
Any question about SAS should be addressed to the XMM-Newton help desk.
Updated on:
April 04, 2006