Download 3. The MONS Telescope requirements

The MONS Payload
Requirements Specification
Rømer System Definition Phase 2000/2001
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Date:
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Classification:
MONS/IFA/PL/RS/0001(2)
22.04.2001
Hans Kjeldsen and Tim Bedding
Jørgen Christensen-Dalsgaard
Jørgen Christensen-Dalsgaard
Open: The document is unclassified and there are no
restrictions in circulation.
The MONS Payload: Requirements Specification
The MONS Payload
Requirements Specification
Teoretisk Astrofysik Center, Institut for Fysik og Astronomi, Aarhus Universitet
This document may only be reproduced with permission of TAC/IFA, except within the Rømer project where
any type of reproduction is allowed.
A draft issue of parts of this document exist in the following document: MONS/AAU/MTR/2000/01(1)
MONS/IFA/PL/RS/0001(2)
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The MONS Payload: Requirements Specification
DISTRIBUTION
This document is for internal use by the Rømer project. All Rømer key persons will be able to access the
document via the DSRI webpage: http://www.dsri.dk/roemer/pro.
MONS/IFA/PL/RS/0001(2)
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The MONS Payload: Requirements Specification
Contents
1. SCOPE
5
2. APPLICABLE DOCUMENTS
6
3. THE MONS TELESCOPE REQUIREMENTS
7
3.1 Photometric precision
7
3.2 Filter passbands
9
3.3 Telescope diameter
10
3.4 Attitude control precision
11
3.5 Optical system
3.5.1 Field stop
3.5.2 Scattered light
15
15
15
3.6 Monitoring of neighbouring stars
16
3.7 Detector
3.7.1 Detector temperature
3.7.2 Analogue-to Digital Converter
17
17
17
3.8 Readout/exposure modes
18
3.9 Data processing
3.9.1 Data processing requirements
20
22
4. SUMMARY OF MONS TELESCOPE REQUIREMENTS
23
MONS/IFA/PL/RS/0001(2)
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The MONS Payload: Requirements Specification
1. SCOPE
This document specifies the requirements for the MONS payload on the Rømer satellite. The document has
been prepared by the Theoretical Astrophysics Center and the Institute of Physics and Astronomy at Aarhus
University as a contribution to the Rømer System Definition Phase (The Danish Small Satellite Programme).
MONS/IFA/PL/RS/0001(2)
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The MONS Payload: Requirements Specification
2. APPLICABLE DOCUMENTS
AD1: Rømer Science Mission Specification
MONS/IFA/MIS/RS/0001(1)
AD2: MONS Field Monitor Requirements Specification and Parallel Science
MONS/IFA/PL/RS/0002(1)
AD3: Requirements for the Star Tracker Parallel Science Programme
MONS/IFA/PL/RS/0003(1)
AD4: MONS Payload Electronics Requirement Specification
TERMA # 255503 DT
AD5: MONS Telescope Optical and Mechanical Sub-System System Definition Phase Design Report
MONS/AUS/PL/RP/0001(2)
AD6: MONS Field Monitor System Definition Phase Design Report
MONS/AUS/PL/RP/0002(1)
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3. THE MONS TELESCOPE REQUIREMENTS
This document describes the requirements on the MONS Telescope design in order to meet the overall
scientific goals defined in the Rømer Science Mission Specification (AD1). The telescope design that is
based on these requirements is described in AD5.
Note that the requirements for the MONS Field Monitor are given in AD2 and the requirements for Star
Tracker parallel science are given in AD3. The requirements for the payload electronics are given in AD4.
3.1 Photometric precision
The main scientific data from the MONS payload will be time series photometric measurements at very high
precision. The extremely accurate measurements needed in order to detect the low amplitude oscillations in
the target stars – the goal is to detect oscillations at a level of a few parts per million (ppm) – is a main driver
for the design of the MONS Payload.
In general the most stable type of observations are differential measurements, i.e., the observed data are
calibrated relative to a reference source that is supposed to represent a stable light source or a light source
that contain the same drift as the target source. This leads to the first requirement for the MONS payload:
The main data from MONS should be differential photometry
There are three ways one can do differential measurements when considering photometry on bright stars:
1. Using a stable artificial light source as a reference.
2. Using stable comparison stars of similar brightness to the target as a reference.
3. Measuring the brightness ratio between two passbands in the target, using the fact that stellar
oscillations show different amplitudes at different wavelengths.
The first option is not preferred because of the problem in finding a sufficiently stable artificial source, and
because light from the calibration source does not traverse the same optical path as the starlight. The
second option is not possible because the science goals of MONS require it to observe some of the brightest
stars in the sky, for which there are no suitable nearby comparison stars. This leads us to choose the third
option for MONS:
The main data from MONS should be two-colour broad-band photometry.
The primary requirement for the mission concerns the photometric precision needed in order to be able to
detect the oscillations. The MONS Science Team have therefore performed a large number of simulations of
expected solar-like oscillations. The simulations consider the following aspects of the MONS mission:

Amplitudes of oscillations based on up-to-date models of excitation of solar-like oscillations in stars
of different masses, temperatures and luminosities.

Mode re-excitation, damping and life time issues. We know (from the Sun) that solar-like oscillations
are only coherent over time scales of up to about one week. We include those effects in the
simulations.

Frequencies of oscillations are calculated from stellar evolution models.

Included in the models are also granulation noise, photon noise and imperfect data sampling due to
the selected orbit (Molniya).
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The simulations indicate that MONS will have no problems in detecting solar-like oscillations, assuming that
the instrumental noise and photon noise will be kept low enough to not being the dominating noise source.
In Figure 3.1.1 we show one example of a simulation, for the target α Centauri A+B (a close double star),
which is one of our primary targets. The simulations show the expected oscillation spectrum after 30 days of
observations in the red-blue colour ratio. Note the rising background at low frequencies, which is due to
granulation noise in the target itself. This granulation is an intrinsic stellar noise source, and we should
design our experiment so that instrumental noise sources fall below it.
Figure 3.1.1. The amplitude spectrum of α Centauri A+B (in colour) after 30 days
of observations with the MONS Telescope.
Based on the simulations we specify:
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The MONS Payload: Requirements Specification
The MONS payload requires high-stability on short time scales and progressively
less stability on longer timescales, to match the shape of the intrinsic stellar
granulation noise.
For periods in the range of 5 minutes (frequencies around 3 mHz), we require that the mean amplitude
A
of the noise be below the stellar granulation noise (for the brightest stars V < 3). The stellar granulation is at
a level of about 0.35 ppm (in colour) after 30 days observing (assuming an observing duty cycle of 85%).
This can be transformed to a requirement on the white noise in the time series using the relation:
A  t   /T ,
where
t
is the rms scatter in the time series and T is the total length of the observing period.
Since T  85%  30d  86400 sec/ d  2.2  10 sec , we get the requirement that
6
 t  295 ppm / sec .
This requires the detection of at least 23 million photons per second in both red and blue images.
3.2 Filter passbands
We expect to observe in blue and red light. Based on temperature sensitivity estimates (the amplitude for a
given oscillation scales approximately as 1 /  ), we have used a solar-type spectrum (G2V) to maximize the
S/N of the oscillations for two-colour broad band photometry. The filters are specified as:



BLUE:
400 – 520 nm
RED:
620 – 780 nm
90 % peak transmission
Figure 3.2.1. The two MONS filters and their position in wavelength. The blue filter
will be centred at 460 nm and the red filter will be centred at 700 nm.
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The MONS Payload: Requirements Specification
Number of photons per pixel per nm per exposure for a G2V star at magnitude V=0
and filter transmission curves as shown above.
We detect the red and blue light on separate areas in the out-of-focus image at the detector. In order to
optimise the S/N, the division will be:


40 % RED image
60 % BLUE image
In practice, there will be two red and two blue sectors.
3.3 Telescope diameter
We need to be able to detect 23 million photons per second in each passband for a V=3 star. We assume
93% reflectivity for the primary mirror and adopt a detector sensitivity as a function of wavelength appropriate
for the baselined CCD (see below).
This requires a telescope aperture of at least 500 to 600 cm2, which corresponds to a diameter of 30 to 33
cm, depending on the diameter of the central obstruction.
We specify 32 cm as the Telescope diameter and about 17 cm as the diameter of
the central obstruction.
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3.4 Attitude control precision
Spacecraft jitter (attitude movements) will have a dramatic effect on the photometric precision if one does not
design the instrument, telescope and platform carefully. This is because there will be small variations on the
sensitivity of the CCD from pixel to pixel and even within each pixel. Most of these variations can be
calibrated with flat-field measurements, but residual mis-calibrations are inevitable.
In order to simulate the effects of photometric noise induced by the platform jitter we have done detailed
computer simulations. The simulations included the effects of flat-field structure (including sub-pixel
variations), ACS jitter and the shape of the telescope PSF (including off-axis aberrations). We also included
readout and photon noise. The assumed flat-field structure is shown in Figures 3.4.1 and the assumed PSF
is shown in Figure 3.4.2.
A key issue is the photometric algorithm used for the on-board data reduction. We measure the total intensity
in a number of fixed annular apertures for each stellar image. The apertures are fixed on the CCD coordinate
system, resulting in practically no sensitivity to flat field errors.
On the ground we expect to do additional processing to correct for the second-order dependence of the
position coming from any small errors in the description of the flat parts of the out-of-focus image as well as
the correction of off-axis aberrations. The idea is to use the data to construct a model for brightness “error”
as a function of position. We have used simple 2-dimensional models to correct for the effect of the secondorder attitude error on the brightness.
The simulations showed that the photometric errors arising from ACS errors form a non-white noise source
whose power spectrum has the same shape as the ACS errors themselves. Following the discussion in
Section 3.1 above, we therefore need to specify our requirements on ACS precision in terms of its power
spectrum. It is not sufficient to specify only the RMS attitude precision, since small variations on fast
timescales (during individual 2-second exposures) are much more serious than large variations on longer
timescales.
We assume that the ACS power spectrum is flat at frequencies below 10 mHz, which should be the case if
the control loop is operating correctly. We then assume that the ACS power spectrum falls off as frequency
squared (i.e., as 1/f in amplitude), as seems likely. The spectrum can then level out at frequencies higher
than 10 Hz.
The ACS power spectrum should be flat below 10 mHz and drop between 10 mHz
to 10 Hz by a factor of more than 100 in amplitude (10000 in power).
If the ACS power spectrum shape is significantly different from that assumed here then further simulations
will be needed to specify new requirements. However, discussions with the Rømer ACS group indicate that
our assumed power spectrum shape is feasible. We have also learned from a preliminary ACS study by the
Rømer ACS group that it is feasible to reach 1.2 arcmin RMS for the accuracy of the source pointing. In the
simulations described below we have therefore used a simple model for the attitude, which is illustrated in
the Figures 3.4.3 onwards.
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Figure 3.4.1: The 1024 x 1024 pixel flat field that is used during our simulations.
The peak-to-peak variations in CCD quantum efficiency are about 3 %. This
corresponds to local RMS scatter of 0.5 % in the sensitivity.
Figure 3.4.2: In order to simulate correctly the ACS effect on the photometric noise
we included a description of off-axis aberrations in the telescope. The aberrations
shown here corresponds to a 5 arcmin off-axis out-of-focus PSF for the optical
system described in Section 3.5.
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Figure 3.4.3: Simulated attitude movements in pitch and yaw for a 1500-second
period. In the left panel we show 30,000 time steps (20 points per sec). At right
only 2 points per second are shown, together with bold lines showing the attitude
movement during 10 different (and widely separated) 2-second exposures.
Figure 3.4.4: Amplitude spectrum of the above time series, normalized to have unit
amplitude at 0.5 Hz. The amplitude is 20-40 units at 10 mHz and 0.1-0.2 units at
10 Hz.
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The MONS Payload: Requirements Specification
Figure 3.4.5: Pitch variations in the 1500-second simulated time series, shown at
full resolution (20 points per sec). The RMS in pitch is 1.2 arcmin.
Figure 3.4.6: Same as above, but showing mean values for each 1-second period.
The RMS of this smoothed time series is 1.19 arcmin.
Figure 3.4.7: Point-to-point variation of the time series in the previous figure (i.e.,
differences between adjacent 1-second averages). The RMS for this time series is
0.255 arcmin.
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The MONS Payload: Requirements Specification
Based on the optical system described below and the following assumptions:



Flat field peak-to-peak:
Local Flat field RMS:
ACS-power at 10 Hz:
3%
0.5 % (pix-to-pix)
below 0.01 % of the ACS-power at 10 mHz
we find that in order to keep the ACS-related noise below 0.1 ppm in amplitude at a frequency of 3 mHz after
30 days observing one should limit the peak-to-peak ACS movements below 40-50 pixels. The total RMS of
the ACS movement should therefore be less than 12-14 pixels. The error scales as the square of the size of
the RMS of the movements in pitch and yaw so in order to allow for a factor of two error in the noise
estimate, we specify a maximum of 13 / 2  9 pixels for the attitude movement. Since the pixel size is
0.013 mm we could also write this requirement as:
ACS should be pointing to better than 0.117 mm (RMS) at the detector plane.
We know from the early ACS studies that the ACS will probably be able to deliver 1.2 arcmin (RMS) or 2.4
arcmin (95 % probability).
3.5 Optical system
Combining the expected ACS performance with the limits on detector-plane movements, we get the
requirement that the image scale at the detector should be at least 10.3 arcmin/mm. This sets a requirement
on the optics:
The telescope focal length should be shorter than 330 mm.
To avoid saturation, the image must be spread over many pixels. With the telescope area specified above,
the CCD will detect 300 million electrons/second in the BLUE and 375 million in the RED for a G2V star with
V=0. We assume the CCD saturates at 35,000 e/pixel. The approximate size of the out-of-focus image at
the detector plane should then be:

image diameter:
9.0 mm (692 pixels)
The baseline design is a telescope with a paraboloidal primary mirror, a field stop at the focus and the CCD
placed behind the focus.
To reduce photometric errors from mis-calibration of the flat-field response to acceptable levels, the
simulations show that the PSF should be uniform to within 1 or 2% over the scale of a few pixels.
3.5.1 Field stop
Since the detector is not in focus, a field stop is needed to reject light coming from the field around the target
star, which will contain neighbouring stars and background zodiacal light. The field stop also serves to reject
scattered light (see below). The diameter of the field stop should be as small as possible, while being large
enough to ensure that the target star is not vignetted during the maximum expected attitude
errors (see AD5).
3.5.2 Scattered light
The need to reject scattered light sets requirements on the baffling and on the quality of the mirror surface.
We require that the total amount of scattered light reaching the detector from all sources (Sun, Earth, Moon,
bright stars and planets) be lower than the zodiacal background by a factor of three. Note that zodiacal light
dominates CCD readout noise as the main source of background noise (see below). The total amount of
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The MONS Payload: Requirements Specification
6
zodiacal light reaching the detector is around 10 photons/s (combined blue and red). We therefore specify
that scattered light should result in less than 3  10 photons/s reaching the detector.
5
Using the fact that a star of magnitude V=0 will produce 7  10 photons/s at the detector, we reach the
following requirements for stray light rejection in the MONS Telescope:
8
Required rejection factor:
The Sun:
V = -26.8
The Earth:
V = -21
The Moon:
V > -12.7
Jupiter:
V = -2.7
19
Flux: 3.7  10 photons/sec
17
Flux: 1.8  10 photons/sec
13
Flux: 8.4  10 photons/sec
9
Flux: 8.4  10 photons/sec
8  10 15
2  10 12
4  10 9
4  10 5
For Jupiter and any other bright planets or stars, this rejection factor should easily be achieved. For the
Earth and Moon, it seems clear that no direct light should be allowed to fall on the primary mirror, which in
turn sets a requirement on the minimum length of the telescope tube. For scientific reasons (to allow access
to all primary targets during the mission), we will need to observe down to 30 degrees from the limb of the
Earth. With a 32 cm telescope we therefore need a tube length of at least 48 cm.
For the Sun, the required rejection level probably means that sunlight should not strike the inside of the
telescope tube, which sets a requirement for a sunshade. In the current design, sun shading will be provided
by a deployable cover, which also protects the telescope during integration and launch.
3.6 Monitoring of neighbouring stars
Light from faint neighbouring stars that lie close to the target star will pass through the field stop and be
detected. Some of these stars may be variable at much higher levels than the target, giving signals in the
total light level (and colour ratio) that would be comparable to the oscillation signal from the target itself. For
example, a pulsating star that is 1000 times (7.5 magnitudes) fainter than the target and variable at the level
of 1000 ppm (0.1%) will produce a 1 ppm variation in the total light (about 0.6 ppm in colour). Many stars are
variable at this level. Without detailed knowledge of the variability or otherwise of all neighbouring stars, one
would never know for sure whether unusual peaks in the power spectrum (e.g., at low frequencies) were due
to contamination. This presents an unacceptable risk to the primary science mission.
We require that all neighbouring stars be monitored for variability, to determine all
oscillations with periods and amplitudes that would contaminate the target power
spectrum.
Monitoring should be done by differential photometry of the field surrounding the target, with observations
made over the same length of time as those with the MONS Telescope and preferably simultaneously. We
require that any variability in a neighbouring star which produces a peak in the final MONS Telescope
amplitude spectrum more than 2.5 times the noise level should be detected at the 4-sigma level by this
photometric monitoring. For example, if the target has magnitude V=3 then a neighbouring star at V=10.5
(1000 times fainter) must be measured to a precision of 800 ppm (0.08%).
Achieving such precision and temporal coverage from ground-based observations is not feasible. We
therefore require that a Field Monitor be included in the MONS payload. The requirement specifications for
the MONS Field Monitor, together with a discussion of parallel science that would be possible, are given in
AD2.
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3.7 Detector
High-precision CCD photometry is a well-documented technique. However, we should not just base our
photometric requirements on the knowledge collected from the last 10 years of ground-based observing, but
should optimise the detector specifications on our special needs. In the simulations described in this
document, we have taken a number of CCD properties into account, including spatial sensitivity variations (at
pixel and sub-pixel levels), read-out properties (linearity, dark current, read noise), radiation damage from
cosmic rays (mainly particle radiation) and temporal sensitivity variations (drifts in the detector performance).
3.7.1 Detector temperature
Although the CCD will be protected by aluminium shielding, we have to consider the damage to the CCD
chip from cosmic rays (high energy protons and particle radiation). The radiation levels will be significantly
higher than is normally the case even for the CCD chips in the Hubble Space Telescope focal plane
instruments. Cosmic rays will result in the following damage to the CCD or the photometric signal:

When a cosmic ray hits the chip a few pixels will experience a higher count level (it may in some
cases even saturate the detector in the pixels which are hit by the cosmic high energy particle).

The dark current may increase (in the long run it may change significantly). The worst effect may not
be the dark current level, since we expect to use only short exposures. A type of dark noise called
random telegraphic pixel noise (RTS noise) may dominate the dark noise contribution. This is
caused by dark current flipping between two discrete values for a given pixel, which will result in a
non-white noise. The only way to limit this noise source is by having a low level for the dark noise.

Charge Transfer Efficiency (CTE) will decrease with time (due to traps). If the CTE becomes too low
we may see trailed images as the CCD is read out.

Hot pixels are pixels with an extremely high dark current (it could in principle saturate a single pixel).
The number of those hot pixels will depend on the CCD temperature and the level of radiation.

Traps (defects that absorb and release charge as it is clocked through the defect area) will be
generated by the strong radiation. A low temperature will limit the negative effect of those traps.
The effects mentioned above indicate that a low temperature for the CCD chip is preferred. This conclusion
is supported by the ESA studies behind the two asteroseismic projects (PRISMA and STARS) proposed for
ESA medium-size missions (M2 in 1993 and M3 in 1996), and the resent ESA F2/F3 mission proposal,
EDDINGTON, which has now been chosen by ESA Space Science Advisory Committee (SSAC) as a
reserve mission to fly before 2010. Also a detailed study by Marconi Applied Technologies on the radiation
damage performance of their CCDs (S&C906/424, 17/2/2000) indicate that one should gain at a low
operation temperature for the main instrument – MONS CAM. CCD temperature stability studies by Niels
Bassler (TERMA report, 2000) indicate that the sensitivity to variable chip temperature should not be
ignored. The same study also indicated that the sensitivity to thermal instabilities will be low enough to
remove if one stabilize the operation temperature of the CCD.
During the Detailed Design Phase we plan to perform test of CCDs, their temperature response as well as
their sensitivity to radiation. At present we use a set of safe requirements for the CCD temperature.
Temperature requirements:
Operation Temperature:
Temperature Stability:
Temperature measured to (HK):
-70 degC to –100 degC
0.1 degC (95 % the of time)
+/- 0.02 degC (95 % the of time)
3.7.2 Analogue-to Digital Converter
A 12-bit ADC is needed to fully sample the dynamical range of the signal.
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3.8 Readout/exposure modes
In order not to saturate the CCD for the brightest stars and in order not to be dominated by read out noise
during normal operation we need to define three different types of read-out modes that will depend on the
brightness of the target star. The three modes can be seen in the following table. Mode I is for the brightest
stars, mode II for medium brightness stars and mode III for faintest stars:
I
V < 1.5
Exp.time:
Pixel-binning (on chip):
#pix (from MONS CAM CHU):
Readout time:
Software binning:
Frame size:
V = 0, Sp=G2V:
Total:
R.O.N:
Zodiacal background:
Noise: V=0, G2V:
V=1.5, G2V:
II
1.5 < V < 3.0
Exp.time:
Pixel-binning (on chip):
#pix (from MONS CAM CHU):
Readout time:
Software binning:
Frame size:
V = 1.5, Sp=G2V:
Total:
R.O.N:
Zodiacal background:
Noise: V=1.5, G2V:
V=3, G2V:
III
V > 3.0
Exp.time:
Pixel-binning (on chip):
#pix (from MONS CAM CHU):
Readout time:
Frame size:
V = 3, Sp=G2V:
Total:
R.O.N:
Zodiacal background:
Noise: V=3, G2V:
V=5, G2V:
2 sec
2 x 2 pixels
262.1 k
561 msec
4 x 4 pixels
128 x 128 pix.
14900 e/pix (2x2bin) BLUE
27800 e/pix (2x2bin) RED
6.02 e+08 e/exp BLUE
7.48 e+08 e/exp RED
4020 e/exp BLUE
3280 e/exp RED
9.55 e+05 e/exp BLUE
8.22 e+05 e/exp RED
55 ppm/exp (photon: 55)
114 ppm/exp (photon: 109)
2 sec
4 x 4 pixels
65.5 k
439 msec
2 x 2 pixels
128 x 128 pix.
14900 e/pix (4x4bin) BLUE
27800 e/pix (4x4bin) RED
1.51 e+08 e/exp BLUE
1.87 e+08 e/exp RED
2010 e/exp BLUE
1640 e/exp RED
9.55 e+05 e/exp BLUE
8.22 e+05 e/exp RED
111 ppm/exp (photon: 109)
230 ppm/exp (photon: 219)
2 sec
8 x 8 pixels
16.4 k
409 msec
128 x 128 pix.
14900 e/pix (8x8bin) BLUE
27800 e/pix (8x8bin) RED
3.76 e+07 e/exp BLUE
4.68 e+07 e/exp RED
1005 e/exp BLUE
820 e/exp RED
9.55 e+05 e/exp BLUE
8.22 e+05 e/exp RED
223 ppm/exp (photon: 219)
614 ppm/exp (photon: 550)
As can be seen from the table above, for all stars brighter than magnitude V=5 we are dominated by photon
noise. The noise is also illustrated in the following figure.
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The MONS Payload: Requirements Specification
The rms scatter per exposure (2 sec) in the colour ratio for stars of different
magnitude. The CCD is operated in three different readout modes depending on
the magnitude of the star. Mode I is for the brightest stars, mode II for the medium
brightness stars and mode III for the faintest stars. To get to 4-sigma detection
sensitivity in the amplitude spectrum after 30 days (duty cycle: 85 %) one should
divide the scatter with 150.
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3.9 Data processing
Each 2-second exposure from the MONS Telescope produces a 128 x 128 pixel binned image. For
debugging purposes, it should be possible to download entire images. During normal operations, however,
the images will be processed on-board.
The images will be processed in blocks of 5, as illustrated in Figure 3.9.1 and described in the following
sections. The main reason for grouping into blocks of 5 is to allow removal of cosmic rays by median
filtering, as described below.
The top row of Figure 3.9.1 shows 5 consecutive 2-second images. Some cosmic ray hits are visible, and
each image has a brightness profile due to the telescope PSF. The first step is to divide by the flat-field. We
then measure the centroid of each image, which reflects the average pointing offset during the exposure.
This knowledge is used to calculate the PSF appropriate to that offset and divide by it. We also subtract the
background, which is measured in regions of the CCD that are not illuminated by the star. The second row
of the Figure shows the images after these corrections have been made. Note that the images are now
much flatter, although the cosmic rays still remain.
The third row of Figure 3.9.1 is calculated by subtracting the perfectly corrected on-axis image from each
image in the second row. Note that these images are shown here for illustration and would not be calculated
on-board (the perfectly corrected image would not be known!). These residuals show that the PSF
correction was not perfect, due to factors such as small errors in the centroiding, an imperfect model for the
PSF and/or telescope pointing changes during the exposure. But the deviations from flatness which they
illustrate must be corrected before the five images can be median-filtered. Since the deviations basically
appear as a linear slope across the images, we can use the four-fold reflection symmetry of the images to
remove the deviations. That is, for each pixel in an image we identify the three other pixels that are reached
via reflections in the horizontal or vertical axes (or both). See first image in row 2 for an example. Those four
pixels are then all replaced with their mean.
The fourth row of the Figure shows the results. Note that these images are four-fold redundant (as can be
seen by the four-fold appearance of cosmic rays). The fifth row again shows the residuals of the five images
with respect to the perfectly corrected image. These residual images, again shown only for illustration,
demonstrate that the images in row four are now extremely flat.
We are now ready to combine the five images in row four using a median filter. We keep track of the relative
contributions of the five images to the median, since if these are all equal we can be confident there is no
significant drift in total flux (as might be caused by variations in CCD gain or timing). The final image,
together with its deviation from the perfectly corrected image, are shown in Figure 3.9.2. Note the successful
removal of cosmic rays.
Photometry can now be performed on the red and blue sectors of this image. This is done by counting
photons within apertures that are fixed in CCD coordinates. This is to ensure that the same CCD pixels are
used throughout the whole time series, which removes sensitivity to imperfect flat-fielding. In practice, we
will downlink counts for a number of sub-apertures, including some outside the stellar images (to measure
any uncorrected background).
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The MONS Payload: Requirements Specification
Figure 3.9.1. A series of five consecutive 2-second images during various stages
of processing (see text).
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The MONS Payload: Requirements Specification
Figure 3.9.2: Image resulting from combining the five images with a median filter.
3.9.1 Data processing requirements
The image processing described above will require:
RAM buffer
0.295 Mbyte (16-bit RAM)
Each block of five 2-second images will produce 24 numbers (4 red counts, 4 blue counts, 2 background
counts, 5 x/y centroid coordinates and 4 relative contributions to the median), which must be stored and
transmitted to the ground. The amount of data, assuming 4-byte floating numbers, is 9.6 bytes/sec. We will
observe about 20.5 hours per day, producing 708 kbytes/day, giving the requirement
Telemetry downlink
0.71 Mbyte/day
To allow for ground station downtime, we should be able to store the data for 48 hours:
Data storage
0.71 Mbyte (32-bit RAM)
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The MONS Payload: Requirements Specification
4. SUMMARY OF MONS TELESCOPE REQUIREMENTS
Assumed orbit
Molniya
Observing method
two-colour differential photometry
Optical design
on-axis paraboloidal telescope with field stop at the focus and
detector behind the focus
Telescope aperture
320 mm diameter, central obstruction about 170 mm
Telescope focal length
< 330 mm
Image Size
9.0 mm diameter
Image PSF
uniform within 1–2% over small scales (several pixels)
Filter passbands
Red
620–780 nm (FWHM), 40% of aperture
Blue
400–520 nm (FWHM), 60% of aperture
Filter peak transmission
90%
Primary mirror reflectivity
93%
Pointing exclusion
able to point down to 30 degrees from limb of Moon and Earth
Telescope tube length
> 480 mm
ACS precision
see Section 3.4
Field stop
as small as possible, but without vignetting (99% of the time)
Monitoring neighbouring stars requires MONS Field Monitor
Scattered light levels
see Section 3.5.2 for numerical values
Detector
Marconi CCD 47-20 AIMO (back/thinned)
CCD electronics
12 bit ADC
CCD Cooling
–70 to –100 oC.
Stability
0.1oC
Measured to 0.02 oC
CCD readout modes
see Section 3.8
RAM buffer
0.295 Mbyte (16-bit RAM)
Data storage
0.71 Mbyte (32-bit RAM)
Telemetry downlink
0.71 MByte/day
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