Article Index

THESEUS Mission Payload and Profile

The scientific goals which come from a full exploration of the early Universe requires the detection of a factor ten more GRBs (about 100) in the first billion years of the Universe (z > 6), in the 3 years prime mission life time of THESEUS. Such a requirement is well beyond the capabilities of current and near future GRB detectors (Swift/BAT, the most sensitive one, has detected only very few GRBs above z = 6 in 10 years). As supported by intensive simulations performed by us and other works in the literature, the needed substantial increase of high-z GRBs requires both an increase of ~1 order of magnitude in sensitivity and an extension of the detector passband down to the soft X-rays (0.5 – 1 keV). Such capabilities must be provided over a broad field of view (~1 sr) with a source location accuracy < 2’ , in order to allow efficient counterpart detection, on-board spectroscopy and redshift measurement and optical and IR follow-up observations.

Such performances can best be obtained by including in the payload a monitor based on the lobster-eye telescope technology, capable of focusing soft X-rays in the 0.3 – 6 keV energy band over a large FOV. Such instrumentation has been is under development for several years at the University of Leicester, has a high TRL level (e.g., BepiColombo) and can perform all-sky monitoring in the soft X-rays with an unprecedented combination of FOV, source location accuracy (<1-2’) and sensitivity, thus addressing both main science goals of the mission. An onboard infrared telescope of the 0.5-1m class is also needed, together with spacecraft fast slewing capability (e.g., 5-10°/min), in order to provide prompt identification of the GRB optical/IR counterpart, refinement of the position down to ~arcsec precision (thus enabling follow-up with the largest ground and space observatories), on-board redshift determination and spectroscopy of the counterpart and of the host galaxy. The telescope may also be used for multiple observatory and survey science goals. Finally, the inclusion in the payload of a broad field of view hard X-ray detection system covering the same monitoring FOV as the lobster-eye telescopes and extending the energy band from few keV up to several MeV will increase significantly the capabilities of the mission. As the lobster-eye telescopes can be triggered by several classes of transient phenomena (e.g., flare stars, X-ray bursts, etc), the hard X-ray detection system provides an efficient means to identify true GRBs and detect other transient sources (e.g., short GRBs). The joint data from the three instruments will characterize transients in terms of luminosity, spectra and timing properties over a broad energy band, thus getting fundamental insights into their physics. In summary, the foreseen payload of THESEUS includes the following instrumentation:

  • Soft X-ray Imager (SXI, 0.3 – 6 keV): a set of 4 lobster-eye telescopes units, covering a total FOV of ~1sr with source location accuracy < 1-2’;
  • InfraRed Telescope (IRT, 0.7 – 1.8 μm): a 0.7m class IR telescope with 10’x10’ FOV, for fast response, with both imaging and spectroscopy capabilities;
  • X-Gamma rays Imaging Spectrometer (XGIS, 2 keV – 20 MeV): a set of coded-mask cameras using monolithic X-gamma rays detectors based on bars of Silicon diodes coupled with CsI crystal scintillator, granting a ~1.5sr FOV, a source location accuracy of ~5 arcmin in 2-30 keV and an unprecedently broad energy band.

The mission profile should include: an onboard data handling units (DHUs) system capable of detecting, identifying and localizing likely transients in the SXI and XGIS FOV; the capability of promptly (within a few tens of seconds at most) transmitting to ground the trigger time and position of GRBs (and other transients of interest); and a spacecraft slewing capability of ~10-20°/min). The baseline launcher / orbit configuration is a launch with Vega-C to a low inclination low Earth orbit (LEO, ~600 km, <5°), which has the unique advantages of granting a low and stable background level in the high-energy instruments, allowing the exploitation of the Earth’s magnetic field for spacecraft fast slewing and facilitating the prompt transmission of transient triggers and positions to the ground.

On-board Instruments

Following the scientific requirements described before, The baseline Instrument suite configuration of THEUS payload includes 5 lobster-eye modules (F=300mm), a 70cm class IR telescope and a set of hard 25 X-ray / soft gamma-ray detectors based on Si+CsI(Tl) coupling technology covering the total FOV of the lobster-eye modules. In summary, the scientific payload of THESEUS will be composed by:

  • Soft X-ray Imager (SXI): a set of 4 « Lobster-Eye » X-ray (0.3 - 6 keV) telescopes covering a total FOV of 0.1 sr field with 0.5 – 1 arcmin source location accuracy, provided by a UK led consortium;

  • InfraRed Telescope (IRT): a 70 cm class near-infrared (up to 2 microns) telescope with imaging and moderate spectral capabilities provided by a France led consortium (+ ESA, Swiss and German);

  • X-Gamma-rays Imaging Spectrometer (XGIS): spectrometer based on 4 detection units based on SDD+CsI(Tl) modules (2 keV – 20 MeV) , covering twice the FOV of the SXI. This instrument will be provided by an Italian led consortium (+Spain);

All the Instrument shall be equipped by an: Instrument Data Handling Unit (I-DHU): interfacing each of the three instruments with spacecraft. Provided by a German led consortium (+ Poland).

The Soft X-Ray Imager (SXI)

The THESEUS Soft X-ray Imager (SXI) comprises 4 DU. Each DU is a wide field lobster eye telescope using the optical principle first described by Angel (1979) with an optical bench as shown in Figure 3.1. The optics aperture is formed by an array of 8x8 square pore Micro Channel Plates (MCPs) supplied by PHOTONIS France SAS, (Avenue Roger Roncier, 19100 Brive La Gaillarde, France). The MCPs are 40x40 mm2 and are mounted on a spherical frame with radius of curvature 600 mm (2 times the focal length of 300 mm). The open aperture provided by each plate is 38x38 mm2 so that a 1 mm wide strip around the edges is available for mounting/gluing. Table 1 summarizes SXI characteristics. The mechanical envelope of a SXI module has a square cross-section 320x320 mm2 at the optics end tapering to 200x200 mm2 at detector. The depth of the detector housing is 200 mm giving an overall module length of 500 mm. The optics assembly has a mass of 2.5 kg, the detector assembly 13 kg and the structure 2.5 kg giving an overall module mass of 18 kg.


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Energy band (keV)


Telescope type:

Lobster eye

Optics aperture (mm2)


Optics configuration

8x8 square pore MCPs

MCP size (mm2)


Focal length (mm)


Focal plane shape


Focal plane detectors

CCD array

Size of each CCD (mm2)


Pixel size (µm)


Pixel Number


Number of CCDs


Field of View (square deg)

~1 sr

Angular accuracy (best, worst) (arcsec)

(<10, 105)

Power [W] 27.8
Mass [kg] 40
Figure 1. Optical elements of a SXI module. Table 1. The SXI characteristics.


The left-hand side of Figure 2 shows the optics frame of the breadboard model for the SVOM MXT lobster eye telescope which comprises 21 square MCPs mounted over a 5x5 grid (the corners are unoccupied for this instrument). The front surface is spherical with radius of curvature 2000 mm giving a focal length of 1000 mm. The design proposed for SXI uses the same plate size and exactly the same mounting principle but a shorter focal length, (300 mm), so the radius of curvature of the front surface must be 600 mm. The right-hand panels of Figure 2 shows a schematic of a single plate and a micrograph that reveals the square pore glass structure. The focal plane of each SXI module is a spherical surface of radius of curvature 600 mm situated a distance 300 mm (the focal length) from the optics aperture. The detectors for each module comprise a 2x2 array of large format CCDs baselined to be supplied by e2v (e2v technologies (UK) Ltd, 106 Waterhouse Lane, Chelmsford, Essex, CM1 2QU, England), where each CCD has an active area of 81.2x81.2 mm2 (1/6 to be used as a frame store). The detectors are tilted to approximate to the spherical focal surface.


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Figure 2. Left: The SVOM MXT lobster eye optic aperture frame. Top right: A schematic of a single square pore MCP. Bottom right: A micrograph of a square pore MCP showing the pore structure. This plate has a pore size d=20 microns and a wall thickness w=6 microns.


SXI Calibrations

SXI optic: X-ray beam line testing (University of Leicester or Panter MPE) to measure the focal length, the effective area and the point spread function as a function of off-axis angle and energy. SXI detector: Vacuum test facility (University of Leicester or OU) to measure the gain and energy resolution as a function of energy. SXI end-to-end: X-ray beam line facility (University of Leicester or Panter, MPE) to confirm the focal length and measure the instrument effective area and PSF as a function of photon energy and off-axis angle. SXI in orbit calibration: use cosmic sources to confirm in-flight alignment, plate scale, point spread function, effective area, vignetting and energy resolution. Regular monitoring of cosmic sources to monitor calibration. Optional internal X-ray source above CCDs to measure CTI.


SXI Performance, Sensitivity and Data Rate

The imaging area of the CCDs sets the field of view of each module. Each CCD has dimensions 81.2x81.2 mm2 which, with a focal length of 300 mm, and allowing for the frame store gives an active area of 15.5x12.92 square degrees. The field of view of each DU is provided by 4 CCDs giving 801.3 square degrees. Thus, a compliment of 4 SXI modules has a total field of view of 3200 square degrees (0.9 steradians).


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Figure 3 The point spread function of the SXI.


The point spread function is shown in Figure 3. The inner dotted square shows the off-axis angle at which the cross arms go to zero as determined by the L/d ratio of the pores. For optimum performance at 1 keV we require L/d=50. The outer dotted square indicates the shadowing of the cross-arms introduced by the gap between the individual MCPs in the aperture. The central true-focus spot is illustrated by the projection plot to the left. The FWHM is 4.5 arc minutes and all the true-focus flux is contained by a circular beam of diameter 10 arc minutes. The collecting area, within a 10 arc minutes beam centered on the central focus, as a function of energy is shown in Figure 4. The optics provides the area plotted in black. The red curve includes the quantum efficiency of the CCD and the transmission of the optical blocking filter comprising a 60 nm Aluminium film deposited over the front of the MCPs and 260 nm of Aluminium plus 500 nm of parylene on the surface of the detectors. Because the angular width of the optics MCP-array is 2.3 degrees larger than the CCD-array the field of view is unvignetted at 1 keV and above so the collecting area shown in Figure 5 is constant across the field of view.


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Figure 4. Collecting area as a function of energy. The black is the optics only. The red curve includes the quantum efficiency of the CCD and the transmission of the optical blocking filter.     Figure 5. The position accuracy of the SXI as a function of source and background count. R90 is the error radius that contains 90% of the derived positions


Using the Rosat All-sky Survey data we can estimate the count rate expected from the diffuse sky (Galactic and Cosmic) and point sources. The average is 1.12x10-5 cts s-1 per square arc minute, although at the Galactic poles it will be a factor of 2 less than this and in the Galactic plane a factor of 3 or so higher. The expected particle background rate in the CCDs is 0.0013 cts s-1 mm-1 so the mean background rate (sky+particles) is 0.0016 cts s-1 in a circular beam of diameter 10 arc minutes. The sensitivity to transient sources using this background rate and a false detection probability of 1.0x10-10 is shown as a function of integration time in the science requirements web-page. For longer integration times the source count required rises, e.g. to 30 counts for a 10000 second integration. The count rate expected from a typical GRB is 2.0 cts s-1. We find that 94% of the Swift BAT bursts (before Sept. 16 2010) would be detected by the SXI. The X-ray light curve of the afterglow would be detected to >1000 seconds after the trigger for a large fraction of the bursts.

Figure 5 shows how the position accuracy R90 (the radius of a circle which contains 90% of the derived positions) improves as the source count S increases. If we include the background count then R90=C.S/(S+gB)1/2 where C=255 arc seconds. The constant g=0.86 depends on the relative size of the focused spot and the beam used. At the sensitivity limit we get R90=105 arc seconds in 50 seconds integration and R90=74.5 arc seconds in 2000 seconds integration. Of course when source are detected well above the threshold the R90 will be much smaller. If we get >1000 source counts then R90<10 arc seconds and the position accuracy will be limited by the systematic errors in the aspect solution. The average count rate per module is ~40 cts s-1 dominated by the diffuse X-ray sky component. This will rise to ~120 cts s-1 in the Galactic plane. Bright transients will produce 10 cts s-1. Therefore the total maximum count rate expected from 5 modules is ~650 cts s-1 and the average is ~250 cts s-1.


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Figure 6. The two stage trigger algorithm. Top left-hand panel: the detected event distribution ΔT=4 seconds and a source count rate of 40 cts s-1 over the full PSF. The cross-arms are rotated wrt the detector axes to demonstrate how this can be handled. The Top right-hand panel: the detected event distribution in the patch of sky aligned to the cross-arm axes of the PSF (shown as the red rectangle in the top left-hand panel). The red cross-patch indicates the area used for the second stage of the algorithm. Bottom panels: the histograms along columns and rows in the patch.


Trigger Algorithm

The ideal algorithm would be some form of matched filtering using the full PSF distribution but because of the extent of the PSF this would be far too computationally heavy. At the other extreme a simple scheme would be to search for significant peaks using the cell size commensurate with the central peak in the PSF. This would be faster but utilizes only ~25% of the total flux detected. The scheme described below is a two stage process which exploits the cross-arm geometry but avoids computationally expensive 2-D cross-correlation. For the 1st stage the focal plane is divided into square patches with angular side length ~4d/L=1/12 radians aligned with the cross-arm axes. The dotted central square shown in Figure 3 indicates the size and orientation of such a patch. The optimum size of such patches depends on the HEW of the lobster-eye optic and the background count rate. The patches could correspond to detector elements or tiles in the focal plane, e.g. CCD arrays. The peak profile is the line spread profile of the central spot and cross-arms of the PSF. The remainder of the histogram distribution arises from events in the cross-arm parallel to the histogram direction, the diffuse component of the PSF and any diffuse background events not associated with the source. Because we are looking for transient sources the fixed pattern of the steady sources in the field of view at the time would have to be subtracted from the histogram distributions. As the pointing changes this fixed pattern background would have to be updated. A transient source is detected if a significant peak is seen in both histograms.

The sensitivity of detection and accuracy of the derived position of the source within the patch depends on the bin size of the histograms, the HEW of the central peak of the PSF and the background. For the most sensitive detection the bin size should be approximately equal to the HEW but this will limit the accuracy of the position. If the bin size of the histograms is chosen to be significantly smaller than the width of the HEW then the histograms can be smoothed by cross-correlation with the expected line width profile of the peak-cross-arm combination. Using the smaller bin size the histograms can also be used to estimate the position centroid of the source within the patch. The significance used for this first stage should be low, e.g. 2.5 sigma. This will provide candidate positions for the second stage.

For each of the candidate positions identified in the first stage a cross-arm patch is set up to cover just the detector area which is expected to contain a fraction of the full cross-arms and the central peak in the PSF. The cross-patch dimensions are changed depending on the integration time ΔT. For short integration times the total background count will be small and the cross-patch size is set large to capture a large fraction of the counts from the cross-arms and central peak.

The above considers a single value of ΔT. We envisage that a series of searches would be run in parallel each using a different integration time so that the sensitivity limit as a function of ΔT is covered. The basic scheme is illustrated in Figure 6. The total source count assumed for this illustration was 40 spread over the full PSF as plotted in Figure 3. The bin size used for the histograms was 1 arc min, and the HEW of the central peak of the PSF is approximately 4 arc mins. We have tested the algorithm over a range of integration times and background conditions. It achieves the sensitivity limit plotted in the science requirements web-page. When a significant transient peak is identified the position must be converted to sky coordinates using the current aspect solution (from Payload star trackers). Positions of all transients found must be cross-correlated with known source catalogues, e.g. Rosat All-Sky Survey, Flare stars, Swift BAT catalogue etc. Any position which does not match a known position must be passed to the Space Craft as a potential trigger position.

The processing required to implement the above is as follows:

1) Extract frames from the CCD at ΔT=2. seconds – this is the fastest rate set by the frame time. 2) Apply event reconstruction algorithm to the frames to give an event list with positions in CCD pixels and a pulse height. 3) Convert the pixel positions into a local module coordinate frame which is aligned to the cross-arms of the PSF. 4) Accumulate counts in the 1-D histograms. 5) Subtract the fixed source/background pattern from the histograms. 6) Scan the histograms for significant peaks and extract candidate positions for further analysis. 7) Set up the cross-arm mask at candidate positions to look for significant peaks. Calculate an accurate position in the local module coordinate frame for the peak. 8) Convert this position into global sky coordinates (quaternion). 9) Check positions against on-board catalogues to weed out known sources. 10) Communicate unidentified transients to the Space Craft. 4)-7) must be repeated for different Δt values e.g. 2, 20, 200, 2000 seconds.

The X-Gamma Spectrometer (XGS)

The X-Gamma ray Imaging Spectrometer (XGIS) comprises 3 units (telescopes). The three units are pointed at offset directions in such a way that their FOV partially overlap. Each unit (Figure 7) has imaging capabilities in the low energy band (2 -30 keV) thanks to the combination of an opaque mask superimposed to a position sensitive detector. A passive shield placed on the mechanical structure between the mask and the detector plane will determine the FOV of the XGIS unit for X-rays up to about 150 keV energy. Furthermore the detector plane energy range is extended up to 20 MeV without imaging capabilities. The main performance of an XGIS unit are reported in Table 2.



Energy band

2 keV – 20 MeV

# detection plane modules


# of detector pixel /module


pixel size (= mask element size)

5x5 mm

Low-energy detector (2-30 keV)

Silicon Drift Detector

450 μm thick

High energy detector (> 30 keV)

CsI(Tl) (3 cm thick)

Discrimination Si/CsI(Tl) detection

Pulse shape analysis

Dimension [cm]


Power [W]


Mass [kg]




2 - 30 keV

30 - 150 keV

> 150 keV

Fully coded FOV

9 x 9 deg2



Half sens. FOV

50 x 50 deg2

50 x 50 deg2 (FWHM)


Total FOV

64 x 64 deg2

85 x 85 deg2 (FWZR)

2π sr

Ang. res

25 arcmin



Source location accuracy

~5 arcmin (for >6σ source )



Energy res

200 eV FWHM @ 6 keV

18 % FWHM @ 60 keV

6 % FWHM @ 500 keV

Timing res.

1 μsec

1 μsec

1 μsec

On axis useful area

512 <cm2

1024 cm2

1024 cm2

Table 2. Top: XGS specifications. Bottom: XGIS unit characteristics vs energy range    


The detection plane of each unit is made of 4 detector modules each one about 195x195x50 mm in size detecting X and gamma rays in the range 2 keV – 20 MeV.. For each energy loss in the module, whatever procured by EM radiation or ionizing particle, the energy released, the 3 spatial coordinates and the of the interaction and time of occurrence will be recorded. The basic element of a module (Figure 8) is a bar made of scintillating crystal 5x5x30 mm3 in size. Each extreme of the bar is covered with a Photo Diode (PD) for the read-out of the scintillation light, while the other sides of the bar are wrapped with a light reflecting material convoying the scintillation light towards the PDs.


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Figure 7. Sketch of the XGIS Unit (top left). Sketch of the XGIS module (top right and bottom left). A module (bottom left) is made of an array of 32x32 scintillator bars with Si PD read-out at both ends (top right). Both PD and scintillators are used as active detectors. The PDs readout electronics consist of an ASIC pre-amp mounted near each PD’s anode while the rest of the processing chain is placed at the module sides and bottom. Bottom right: Principle of operation of the XGS detection units: low-energy X-rays in teract in Silicon, higher energy photons interact in the scintillator, providing an energy range extended to three orders of magnitude.


The scintillator material is CsI(Tl) peaking its light emission at about 560 nm. The PD is realized with the technique of Silicon Drift Detectors (SDD-PD) (Gatti 1984) with an active area of 5x5 mm2 matching the scintillator cross section. Crystals are tightly packed in an array of 32x32 elements to form the module. SDD and scintillator detect X- and gamma-rays. The operating principle (see Figure 7 bottom right) is the following. The top SDD-PD, facing the X-/gamma-ray entrance window, is operated both as X-ray detector for low energy X-ray photons interacting in Silicon and as a read-out system of the scintillation light resulting from X-/gamma-ray interactions in the scintillator. The bottom SDD-PD at the other extreme of the crystal bar operates only as a read-out system for the scintillations. The discrimination between energy losses in Si and CsI is based on the different shape of charge pulses.

While the electron-hole pair creation from X-ray interaction in Silicon generates a fast signal (about 10-ns rise time), the scintillation light collection is dominated by the fluorescent states de-excitation time [0.68 µs (64%) and 3.34 µs (36%) for CsI(Tl)] and a few µs shaping time is needed in this case to avoid significant ballistic deficit. Pulse shape analysis (PSA) techniques are adopted to discriminate between signals due to energy losses in Si or CsI. The results we obtained for the separation of the energy losses in the case of an 241Am source (Marisaldi et al. 2004) are shown in Figure 8 (left panel). As can be seen from Figure 8 (middle panel), the ratio of the two pulse heights is approximately constant for pulses of common shape and allows discrimination between interactions taking place in Silicon or in the scintillator. For gamma-rays interacting in the scintillator, combining the signals from the two PDs at the extreme of each bar allows to determine the energy and the depth of the interaction inside the crystal (Labanti et al. 2008b).


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Figure 8. Left: bidimensional spectum of a 241Am source showing interaction in Silicon (top line) and in the scintillator (bottom line). The output of two parallel processing chains of different shaping times (fast / slow chain) is plot one against the other. Middle: distribution of the ratio of the two processing chains for events in Silicon and in the scintillator. The large peak separation indicates the optimal discrimination performance. Right: XGIS unit main building bloks.


XGIS building blocks

In the XGIS HW, the main building blocks (see Figure 11 for one XGIS unit) are:

a) the mask and the FOV delimeter,

b) the scintillator detectors

c) the FEE in both its analogue part (with SDD-PD, ASIC1 and ASIC2), digital part (DFEE) and services (TLM, TLC Power supply).


Silicon Drift Detectors

The SDD-PD design will be derived directly from the Italian REDSOX R&D program funded by INFN, centered on SDD and related electronics studies. Different SDD geometries have been designed, realized and tested. These SDDs are designed at INFN Trieste and produced by Fondazione Bruno Kessler (FBK, Trento, Italy) that since 2013 is capable of processing 6’’ wafers. This kind of technology demonstrated a noise performance of 8.6 e- rms at room temperature (Bertuccio 2014). The SDD-PD design will be tailored for our specific application. An array of 4 SDD-PD has already been produced (see Figure 9) allowing a tight packaging of the scintillating crystals.


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Figure 9. 2x2 SDD array, 5mm side each one designed by INFN Trieste and produced by FBK, Trento.


Front-End Electronics

The Analogue Front End Electronics (AFEE) design will meet the requirements of minimizing the overall noise readout of an SDD taking into account the limited power budget. For the whole XGIS the overall number of SDD-PDs to be read out is 24.576. The functions of the AFEE for a single detector are illustrated in Figure 10 (left panel). Due to the large number of elements, the technology used in the AFEE will be based on ASICs. Certain operations on the signal, as pre-amplification, must be done as close as possible to the SDD-PD to reduce the coupling stray capacitance. As a best solution, two different ASICs will be realized. ASIC-1 dealing with few SDD PDs signals and performing pre-amplification and buffering to transfer the signal to the rest of the processing electronics few tens of cm away, and ASIC-2 dealing with more signals, leading to 512 ASIC-1 and 128 ASIC-2 per module.


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Figure 10. Left: Scheme of the front end electronics we propose. Right - top: VEGA ASICs realised for the readout of 32 SDD channels. Right - bottom: single channel ASIC cell 200x500 μm2 in size.


The ASIC-1 and ASIC-2 design will be based on the VEGA-ASIC design developed at Politecnico di Milano and Università di Pavia for the readout of SDD devices in the REDSOX project frame (Figure 10, right panel) (Ahangarianabhari 2015). The main characteristics of one of the 32 VEGA channels, having a size of 200×500 μm2, are ~20 electrons rms noise when tested with large-area SDDs and 420 μW/channel consumption for the functions of preamplifier, programmable shaper, discriminator and peak stretcher. With this configuration, an energy resolution of 205 eV FWHM at 5.9 keV has been obtained (Campana 2014). ASIC-1s will be placed on the top and bottom of each XGIS module while ASIC-2s will be placed on the lateral side of each XGIS module.

The Digital Front End Electronics (DFEE) interface the FEE with the Instrument Data Handling Unit (I-DHU) and the Payload Power Supply Unit (PSU). Each XGIS module will have its DFEE whose main functions are the following: interface the stretched ASIC-2 output and provide Analog to Digital Conversion (ADC), event time tagging, detector and module identification.

Units Control Electronics (UCE) At XGIS units level (4 modules) the UCE will manage: a) low voltage (LV=3.3 V) and medium voltage (MV =180V) power supply post-regulation and filtering and distribution, b) TM and TLC interface and management.


Data Handling Unit (DHU) and its functions

The whole XGIS background data rate (3 units) towards the DHU is of the order of 6.000 event/sec in the 2-30 keV range and about 3.700 event/sec above 30 keV. Each event received by DHU will be identified with a word of 64 bits (4 for module address, 10 for bar address, 10 for signal amplitude of the fast top channel, 11 for signal amplitude of the slow top channel, 11 for signal amplitude of the slow bottom channel, 18 for time).

DHU functions will be:

  • discriminates between Si and CsI events.

  • For CsI events, evaluates the interaction position inside the bar by weighting the signals of the 2 PD (a few mm resolution expected). Combining this information with the address of the bar (5x5 mm in size) each module becomes a 3D position sensitive detector.

  • Exploiting the 3D capabilities background can be minimized .

  • It continuously calculates along the orbit the event rate of each module in different energy bands (typically 2-30 keV, 30 -200 keV and >200 keV) and on 5 different times scales (eg 1 ms 10 ms, 100 ms 1s 10 s).

  • In the 2-30 keV range and for each unit, it produces images of the FOV in a defined integration time.

  • It holds in a memory buffer all the XGIS data, rates and images of the last 100 (typical) seconds with respect to the current time.

  • Produce maps of the three unit planes with event pixel by pixel histograms in different energy bands (typically 32 with E width varying logarithmically) and with selectable integration times (min 1 sec).


XGIS and the GRB trigger system

XGIS will contribute to the THESEUS’s GRB trigger system in different ways:

a) Qualification of the SXI triggers. The primary role of XGIS is to qualify the SXI triggers as true GRB. The basic algorithm for GRB validation is based on an evaluation of the significance of the count rate variation, calculated as described below. The procedure will be the following:

  1. from the SXI direction given to the event, it is identified one of the three XGIS units in which the event has potentially been detected;

  2. look for an excess of the rates in the modules of this unit in the bands 2-30 keV and 30-200 keV with respect to the average count-rate continuously calculated by DHU.

b) Autonomous XGIS GRB trigger based on data rate. The autonomous GRB trigger for XGIS inherits the experience acquired with the Gamma Ray Burst Monitor (GRBM) aboard BeppoSAX and that acquired with the Mini-calorimeter aboard AGILE (Fuschino et al., 2008) and concerns all the modules of the 3 XGIS units. For each module the above energy intervals (2-30 and 30-200 keV) are considered for the trigger. The mean count rate of each module in each of these bands is continuously evaluated on different time scales (e.g., 10 ms, 100 ms, 1s, 10s ). A trigger condition is satisfied when, in one or both of these energy bands, at least a given fraction (typically ≥3) of detection modules sees a simultaneous excess with a significance level of typically 5 sigma on at least one time scale with respect to the mean count rate.

c) Autonomous XGIS GRB trigger based on images. For each XGIS unit, the 2-30 keV actual images will be confronted with reference images derived averaging n (typically 30) previous images, and a spot emerging from the comparison at a significance level of typically 5 sigma will appear. If one of the above trigger condition is satisfied, event by event data, starting from 100 sec before the trigger are transmitted to ground, the duration of this mode lasting until the counting rate becomes consistent with background level.


Telemetry requirements

XGIS for transient or persistent source observation.

For the study of transient or persistent sources different transmission mode will be selected starting from the photon list and the histogram maps of the units. Typically the TLM load will be maintained below 2 Gbit/orbit transmitting:

  • at low energies (<30 keV) pixel by pixel histograms in various Energy channels (e.g. 32 ch) with variable integration times (e.g. 64 sec).

  • above the 30 keV the whole photon list.

In particular observations (e.g. crowded fields) a photon by photon transmission in the whole Energy range will be selected for a total maximum TM load of 3 Gbit/orbit.

In the case of a GRB trigger all the information available photon by photon is transmitted with a maximum TLM load of 1 Gbit.


XGIS sensitivity

The 5 sigma XGIS 1 s sensitivity with energy in the SXI FOV, is shown in Figure 11, along with the XGIS flux sensitivity vs. observation time at a significance of 5σ in different energy ranges. In Figure 12, the FOV of the XGIS in the 2-30 keV band is compared with the SXI FOV, and the XGIS sensitivity vs. GRB peak energy is compared with that of other instruments.


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Figure 11. Left: XGIS sensitivity vs. energy in 1 second. Right: XGIS sensitivity as a function of exposure time in different energy bands.


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Figure 12. Left: FOV averaged effective area of the XGIS SDD and CsI detecting planes. Right: Sensitivity of the XGIS to GRBs in terms of minimum detectable photon peak flux in 1s (5sigma) in the 1 - 1000 keV energy band as a function of the spectral peak energy (a method proposed by Band 2003). As can be seen, the combination of large effective area and unprecedented large energy band provides a much higher sensitivity w/r to previous (e.g., CGRO/BATSE), present (e.g., Swift/BAT) and next future (e.g., SVOM/ECLAIRS) in the soft energy range, while keeping a very good sensitivity up to the MeV range.

The InfraRed Telescope (IRT)

The InfraRed Telescope (IRT) on board THESEUS is designed in order to identify, localize and study the transients and especially the afterglows of the GRBs detected by the Soft X-ray Imager (SXI) and the X and Gamma Imaging Spectrometer (XGIS).

The telescope (optics and tube assembly) can be made of SiC, a material that has been used in other space missions (such as Gaia, Herschel, Sentinel 2 and SPICA study). Simulations using a 0.7 m aperture Cassegrain space borne NIR telescope (with a 0.23 m secondary mirror and a 10x10 arc min imaging flied of view), using a space qualified Teledyne Hawaii-2RG (2048x2048 pixels) HgCdTe detector (18 μm/pixels, resulting in 0.3 arcsec/pix plate scale) show that for a 20.6 (H, AB) point like source in a for a 300 s integration time one could expect a SNR of ~5 for a point source. The telescope sensitivity is limited by the platform jitter. Taking the latter into account (1 arc sec jitter over 10 s, at 3 σ level), we foresee to limit the image integration time to a maximum of 10 seconds per frame in order to correct for the jitter, and hence such short integration times will induce a high Read-out Noise (RoN, see below) degrading in turn the IRT sensitivity. In addition, due to the APE capability of the platform (2 arc minutes), the high resolution spectroscopy mode cannot make use of a fine slit, and a slit-less mode over a 5x5 arc min area of the detector will be implemented (similarly to what is done for the WFC3 on board the Hubble Space Telescope), with the idea of making use of the rest of the image to locate bright sources in order to correct the frames a posteriori for the telescope jitter. This is possible thanks to the selectable number of outputs (up to 32) of the Hawaii 2RG detector. The same goal could also be obtained by making use of the information provided by payload the high precision star trackers mounted on the IRT.

Hence the maximum limiting resolution that can be achieved by such a system for spectroscopy is limited to R~500 for a sensitivity limit of about 17.5 (H, AB) considering a total integration time of 1800 s. The IRT expected performances are summarized in Table 7.

In order to achieve such performances (i.e. in conditions such that thermal background represents less than 20% of sky background) the telescope needs to be cooled at 240 (±3) K , and this can be achieved by passive means. Concerning the IRT camera, the optics box needs to be cooled to 190 (+/- 5) K and the IR detector itself to 95 (+/-10) K: this allows the detector dark current to be kept at an acceptable level. The Hawaii-2RG provides a high quantum efficiency (80-90%) over a wide energy range, since substrate-removed HgCdTe can simultaneously detect visible and infrared light, enabling spectrographs to use a single focal plane assembly for Visible-IR sensitivity. The maximum wavelength cut-off (50% of quantum efficiency) varies from 1.8 (H band) to 2.5 (K band) μm as a function of operation temperature detector thickness, and Cadmium fraction of the detector.

In the simulations we used a conservative value of 0.1 e-/s/pixel for dark current at 100 K and 10 e- for the RoN At higher temperatures the dark current increases exponentially (already 1.0 e-/s/pixel at 110 K). The cooling of the detector at these low temperatures can hardly be achieved with a passive system in a low Earth orbit such as the one foreseen for THESEUS, due to the irradiation of the radiators of the infrared flux by the Earth atmosphere. A TRL 5 cooling solution for space applications is represented by the use of a Miniature Pulse Tube Cooler (MPTC). Such devices are available e.g. at Air Liquide Advanced Technologies. Air Liquide has studied in detail an engineering model of MPTC for the CNES MICROCARB mission, and is planning to have a full space qualified version (TRL 8) by 2018. Preliminary thermal studies making use of the Earth fluxes computed for the SVOM mission, which will have a similar orbit in terms of altitude but with a higher inclination (30°), indicate that use of such a device coupled to a 0.3 m2 radiator, which periodically faces the Earth atmosphere, would allow to reach the desired temperatures. Due to the tight link between the thermal aspects at optical and instrument level we propose that ESA takes the responsibility delivering the IRT telescope and its associated thermal system (radiators, etc.) including the IRT instrument cooling system (for the camera and the detector).

In order to keep the camera design as simple as possible (i.e. avoiding to implement too many mechanisms, like tip-tilting mirrors, moving slits etc.), we could implement a design with an intermediate focal plane making the interface between the telescope provided by ESA/industry and the IRT instrument provided by the consortium, as shown in the block diagram in Figure 13. The focal plane instrument is composed by a spectral wheel and a filter wheel in which the ZYJH filters, a prism and a volume phase holographic (VPH) grating will be mounted, in order to provide the expected scientific product (imaging, low and high-resolution spectra of GRB afterglows and other transients).


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Figure 13. The IRT Telescope sketch (left) and optical diagram (right).


Specifications of the entire system are given in Table 3. The mechanical envelope of IRT is a cylinder with 80 cm diameter and 180 cm height. A sun-shield is placed on top of the telescope baffle for IRT straylight protection. The thermal hardware is composed by a pulse tube cooling the Detector and FEE electronics and a set of thermal straps extracting the heat from the electronic boxes and camera optics coupled to a radiator located on the spacecraft structure. The overall telescope mass is 112.6 kg and the total power supply is 95W.


Telescope type:


Primary & Secondary size:

700 mm & 230 mm


SiC (for both optics and optical tube assembly)

Detector type:

Teledyne Hawaii-2RG 2048 x 2048 pixels (18 μm each)

Imaging plate scale


Field of view:

10’ x 10’

10’ x 10’

5’ x 5’

Resolution (λ/Δλ):

2-3 (imaging)

20 (low-res)

500 (high-res), goal 1000

Sensitivity (AB mag):

H = 20.6 (300s)

H = 18.5 (300s)

H = 17.5 (1800s)




VPH grating

Wavelength range (μm):

0.7-1.8 (imaging)

0.7-1.8 (low-res)

0.7-1.8 (high-res, TBC)

Total envelope size (mm):

800 Ø x 1800

Power (W):

115 (50 W for thermal control)

Mass (kg):


Table 3. IRT specifications


IRT Observing sequence

1) The IRT will observe the GRB error box in imaging mode as soon as the satellite is stabilized within 1 arc sec. Three initial frames in the ZJH-bands will be taken (10s each, goal 19 AB 5 σ sensitivity limit in H) to establish the astrometry and determine the detected sources colours.

2) IRT will enter the spectroscopy mode (Low Resolution Spectra, LRS) for a total integration time of 5 minutes (expected 5 σ sensitivity limit in H 18.5 (AB)).

3) Sources with peculiar colours and/or variability (such as GRB afterglows) should have been pinpointed while the low-res spectra were obtained and IRT will take a deeper (20 mag sensitivity limit (AB)) H-band image for a total of 60s. These images will be then added/subtracted on board in order to identify bright variable sources with one of them possibly matching one of the peculiar colour ones. NIR catalogues will also be used in order to exclude known sources from the GRB candidates.

  1. In case a peculiar colour source or/and bright (< 17.5 H (AB)) variable source is found in the imaging step, the IRT computes its redshift (a numerical value if 5<z<10 or an upper limit z<5) from the low resolution spectra obtained at point 1) and determines its position. Both the position and redshift estimate will be sent to ground for follow-up observations. The derived position will then be used in order to ask the satellite to slew to it so that the source is places in the in the high resolution part of the detector plane (see below) where the slit-less high resolution mode spectra are acquired. Following the slew, the IRT enters the High Resolution Spectra (HRS) mode where it shall acquire at least three spectra of the source (for a total exposure time of 1800s) covering the 0.7-1.8 μm range. Then it goes back to imaging mode (H-band) for at least another 1800s (TBC). Note that while acquiring the spectra, continuous imaging is performed on the rest of the detector, see Fig 3.16. This will allow to the on board software to correct the astrometry of the individual frames for satellite drift and jitter and allow a final correct reconstruction of the spectra by limiting the blurring effects.

  2. In case that a faint (> 17.5 H (AB)) variable source is found, IRT computes its redshift from the low resolution spectra, determines its position and sends both information to the ground (as for 3a). In this case IRT does not ask for a slew to the platform and stays in imaging mode for a 3600s time interval to establish the GRB photometric light curve (covering any possible flaring) and leading the light curve to be known with an accuracy of <5 %.


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Figure 14. Left: the IRT focal plane division. The blue area (10’x10’) is used for imaging and low resolution spectra. The orange area (5’x5’) is used for high-resolution slit-less spectra. The size of the high-resolution spectral area is limited by the satellite pointing capabilities. Right: Teledyne Hawaii 2RG detector at ESA Payload Technology Validation section during the tests for the NISP instrument on board the EUCLID mission (Credits ESA)


IRT specific calibration needs

The IRT detector single pixel and intra-pixel response will be characterized using a dedicated optical bench in the lab at CEA-Irfu/SAp on the IRT camera qualification model. A similar activity is on-going at CEA in the framework of the EUCLID NISP instrument. The overall instrument calibration can be obtained on ground using the IRT camera flight model connected to a telescope simulator. The most accurate and final calibration of the IRT will be obtained in flight by observing a number of known calibrating sources, and by means of LED lights installed within the camera that uniformly illuminate the detector, as done on the NISP camera on the EUCLID mission.


IRT Telemetry requirements
The telemetry requirements for the different observing modes described above are reported in Table 4. Transmitting full frames is impractical except in the rarest of cases, so in imaging and LRS mode tiles centered on identified sources of interest are chosen for transmission: 64 x 64 pixels in imaging mode, 128 x 64 in LRS mode.  In both modes, tiles around comparison objects are also likely to be necessary.  In HRS mode, a fixed window of 2048 x 64 is used. Clearly, full frame image transfer must be exceptional and tiled mode somewhat restricted. The figures shown are for all observations, so calibration data is included.


Mode and tilesize


Integration (s)

Volume (Mbit/hr)

Full frame (2048 x 2048) [8 MB]


Combined 60


Image tile (64 x 64) [8 KB]


5 sources @ 10s + 10 ref. objs @ 60s


LRS tile (128 x 128) [16 KB]


10 sources @ 60s (6x10s combined)


HRS image (1024 x 1024) ) [2 MB]


1800 (30x60s combined exposures)


 Table 4. IRT data rate.


The Instrument Data Handling Unit (IDHU)

Following the concept behind the organization of the THESEUS instruments as well as the decentralized avionic scheme, each of the three instrument payloads will be equipped with a dedicated Instrument Data Handling Unit (I-DHU) that will serve as their TM/TC and power interface to the spacecraft. The aim of this scheme is to provide sufficient computing power and data storage to the individual instruments and thus to realize a decentralized data handling function.

The mechanical and electrical design of the I-DHUs (described below) will be the same for all three instruments. Also the operating system and basic software that is running on the Processor Board inside each I-DHU will be the same. In addition, an instrument specific data processing software will be run on each I-DHU, implementing e.g. the above mentioned trigger algorithms and event detection codes.

The computational load on the I-DHU is an high margin, which allows the use of a much simpler, off-the-shelf processor with flight heritage on the Processor Board with the load being easily sustainable with still a large margin. The tasks of the I-DHUs can be separated into three main categories, namely data processing, instrument controlling and power distribution. In order to acquire/process the scientific data, the I-DHUs will:

  • collect, process and store the data stream of the respective instrument

  • implement the burst trigger algorithm on SXI and XGIS data

  • implement the IRT burst follow up observation


I-DHU design overview

The I-DHU consists of two main boards that are mounted inside an aluminum case. In addition, each board exists in a cold redundant (identical, non operating) version inside the box, that can replace the nominal board in case of a failure. Switching between the nominal and redundant chain is done from ground via the spacecraft by activating the SpaceWireLink of the respective data processing board.

At the heart of the I-DHU design is the Processor Board. It hosts the central CPU, the mass memory, time synchronization and distribution circuits and the HK/health monitoring acquisition chain. The Processor Board is connected to the spacecraft via one SpaceWire link through which it will receive the telecommands (TC) and send the science and HK/health data. On the other hand, the I-DHU is connected via another Spacewire link to the respective instrument to relay the TCs and acquire the science and HK data. The SpaceWire interface communication is handled directly through the main CPU (description see below), via its existing dedicated hardware interfaces. The CPU is running an RTEMS (Real Time Executive Management System), an operation system (OS) on which the individual software tasks (detailed description below) of this I-DHU will be run in parallel. Dedicated circuitry is foreseen on the Processor Board for the collection, digitization and organization of HK and health data from various voltage, current and temperature sensors. This will reduce the HK tasks on the main CPU, leaving only the science meta data (like rate meters, counters) and dedicated instrument data processing to be done there. The Processor Board will be developed by the IAAT in Tübingen, Germany.

The Power Board within the I-DHU will be developed by the Centrum Badan Kosmicznych, Poland. It will generate the voltages for the Processor Board and distribute the power to the instrument.


I-DHU mechanical and electrical design

The I-DHU design is based on already developed hardware components with flight heritage or qualification and thus a relatively high TRL. It is based on a single mechanical box containing all electronics and connector interfaces. This particular design has a direct heritage from the eROSITA Interface Control Unit boxes and has as such been qualified already with PCBs during the qualification of the eROSITA telescope structural model and has been shown to withstand the conditions of a launch very well. The mechanical and thermal interface to the platform is realized via six fixations on the bottom plate of the box.


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Figure 15. Left: Block diagram of the GR712RC architecture. Right: Overview of I-DHU dimensions



I-DHU Central Processor and Software

The baseline processor for the I-DHU data processing board is the GR712RC (see Figure 15 Left), a dual-core LEON3FT SPARC V8 processor. It has been developed for high reliability Rad-Hard aerospace applications by Gaisler Aerospace. The GR712RC architecture is centered around the AMBA Advanced High-speed Bus (AHB), to which the two LEON3-FT processors and other high-bandwidth units are connected. Low-bandwidth units are connected to the AMBA Advanced Peripheral Bus (APB) which is accessed through an AHB to APB bridge.

The main functions of the I-DHU on-board software are instrument control & health monitoring and science data processing & formatting. The software will be designed in order to allow the instrument to have the complex functionality that it requires to allow itself to be updated and work around problems automatically and with input from the ground. There will be a common software that is the same for all I-DHUs and an extended software part with modules specific for a given instrument. An example of a software module common to all I-DHUs is the determination of the location of a burst or transient event. The time and location of the transient will be transmitted to ground using the onboard VHF system in a < 1 kbit message. The design of the trigger software benefits from the heritage of the SVOM mission concept as well as past team experience on similar systems on BeppoSAX, HETE-2, AGILE as well as the INTEGRAL burst alert system.

I nstrument control will be possible through the software via telecommands from the ground (e.g., power on & off forindividual units, loading parameters for processing/on-board calibration, investigations) and autonomously on-board (e.g., mode switching, FDIR and diagnostic data collection). The software will implement the standard ECSS-style PUS service telecommand packets for housekeeping, memory maintenance, monitoring etc. and some standard telemetry packets for command acceptance, housekeeping, event reporting, memory management, time management, science data etc. The software will be able to send setup information to the instrument and receive and process the housekeeping data coming back and organize these data in a configurable way for a lower rate transmission to the ground. Figure 16 shows the commandable operational modes managed by the I-DHU.


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Figure 16. Overview of the I-DHU operation modes and transitions


I-DHU Resources and temperature

The masses reported in this page are extracted from the CAD-models and part files for the PCBs. The latter are at this point in time only rough estimations and based on values from similar previous designs. The harness assumes a length of 1 m for all cables. Depending on the location of the I-DHU, the values need to be scaled accordingly. The resources to be allocated for each I-DHU are summarized in Table 5 while Table 6 reports the temperature ranges.


Table 5. I-DHU resources allocation summary

Mass [kg]


Volume [mm3]


Power [Watt]



Table 6. I-DHU Temperature Range

Temperature Range

Min. [°C]


Operative Temperature



No-Operative Temperature




The Trigger Broadcasting Unit (TBU)

In case of trigger event it is necessary to provide the trigger data to ground in a short time. It is expected triggered events occurred at the rate of one event per orbit and the data to be sent to ground is very low: the amount is <1 kb/trigger event. To support the prompt transmission of such event data to ground segment it is necessary a link with Earth independent from the satellite TT&C, which can have a link with ground station only once per orbit. The solutions evaluated for an independent and promt burst position broadcast to ground are by the means of:

  • VHF equatorial network (SVOM and HETE-2);

  • Orbcom satellite network (implemented in Agile satellite);

  • Iridium satellite network (tested by INAF/IASF on baloon and ready to be tested in FEES/IOD n-sat);

  • TDRSS link.


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Figure 17. General block diagram of the Trigger Broadcast Unit


The preferred solution is the well proven VHF broacasting based in the utilization of the SVOM equatorial network. Link with SVOM ground segment will be made of 40 ground stations located located around the Earth inside a ± 30°strip.

The satellite to SVOM network link shall be carry out by the Trigger Broadcasting Unit (TBU - see Figure 17) proposed in the baseline as an unit of the Payload Module. The antennas are miniaturised for a better accommodation on PLM and SVM. The VHF ground network is the same of SVOM mission. For the definition of the VHF frequency range a possible choice, according to ITU Article 5, could be 137 – 137.175 MHz, reserved to space research . This band sub-system for trigger transmission is consequently proposed for compatibility with SVOM ground segment. The resources to be allocated for TBU are summarized in Table 7 while Table 8 reports the temperature ranges.


Table 7. TBU resources allocation summary

Mass [kg]


Volume [mm3]


Power [Watt]



Table 8. TBU Temperature Range

Temperature Range

Min. [°C]


Operative Temperature



No-Operative Temperature




THESEUS Mission Configuration

Theseus satellite will operate in a low equatorial orbit (altitude < 600 km, inclination < 5°). This orbital configuration will guarantee a low and stable background level in the high-energy instruments. The mission has been evaluated assuming as baseline a launch with Vega-C.


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Figure 18. Theseus Orbit Configuration (Left) and Theseus Orbit Ground Track (Right).


The Theseus satellite will be equipped with a suite of three instruments: SXI, XGIS and IRT. In summary the Theseus satellite will be capable to:

  • to monitor a large sky sector for detecting, identifying and localizing likely transients/burst in the SXI and XGIS FOV;

  • of promptly (within a few tens of seconds at most) transmitting to ground the trigger time and position of GRBs and other transients of interest;

  • of autonomous (via SXI, XGIS or IRT trigger) orientation in the sky direction of interest;

  • to perform long observation of the sky direction of interest;

Theseus requires a 3-axis stabilized attitude. During its orbital period Theseus will have distinct operational modes:

Survey (burst hunting) mode: relevant to normal operation when SXI and XGIS are searching for transients. The accessible sky for this kind of operation will be determined by the requirement that:

  • Theseus will have a Field of Regards (FoR) defining the fraction of sky which can be monitored of 64°.

  • When monitoring the sky in normal operations, the number of re-pointings per orbit will be of the order of 3, resulting in observations with the SXI and XGIS of about the whole FoR every 3 orbits.

Burst mode: after detecting a GRB or other transient of interest, the satellite is triggered to this mode by SXI and/or XGIS which transmit to satellite computer the quaternion of the area of interest. The satellite will autonomously fast repoint to place the transient within the field of view of the IRT according to the following steps:

  • fast slewing (< ~60°/10min) for IRT LoS pointing to the direction of GRBs and other transient of interest for low resolution spectra,

  • satellite stabilization (RPE) within less than 0.5 arcsec

  • fast data link for GRB coordinate communication to ground within a few min

Follow-up mode: the IRT shall observes inside the full FoV of 10’x10’ the target with the pre-scripted imaging – spectroscopy sequence. In case of IRT high resolution spectra acquisition a further satellite fine slewing (based on IRT source localization) shall be activated to place the IRT LoS inside a reduced FoV of 5’x5’. XGIS and SXI are in specific follow-up data acquisition mode. In this mode the GRB observation shall be performed in a time of 30 minutes. After completion of the transient observation, Theseus will return to Survey Mode to monitor the sky.

IRT observatory mode: the IRT may be used as an observatory for pre-selected targets through a GO programme, driving the pointing of the satellite. XGIS and SXI are observing as in survey mode, with the possibility of triggering the burst mode

No specific orbit parameter change is required during the mission lifetime. Theseus mission can be supported as baseline by a dedicated ground station located in Malindi (3° S, 40° E). Another ground station, located in Alcantara (2°S, 4° W), is supposed as possible back-up in case of Brasilian participation. In conformity with the selected equatorial orbit, both stations will allow highly frequent accesses to the satellite by a contact per orbit.

Contact analysis between the satellite and the ground stations has been performed considering a minimum station elevation angle over local horizon of 10°.

Within the following tables, the average (and the maximum) single contact duration, the number of daily contacts and the resulting average daily contact duration are summarized for every ground station.


Table 9. Accesses to Ground Stations

Ground station

Average contact duration (sec)

Maximum contact duration (sec)

N°contacts /day

Average daily contact duration (sec)




13 to14





13 to 14



Satellite Configuration

The satellite configuration and design take into account a modular approach. The spacecraft platform is divided in two modules, the Payload Module (PLM) and Service Module (SVM). The Payload Module will mechanically support the Instruments DUs (SXI, XGIS and IRT) and will host internally the Instruments ICUs. The instruments SXI and XGIS DUs are accommodated externally on the Payload Module to which they are connected by means of a structural pedestals. The Instrument DUs mechanical fixing to this structure will be designed in order to guarantee thermo-structural decoupling from the rest of satellite.

The Service Module contains all the platform subsystems and provides the mechanical interface with the Launcher. Figure 19 shows the spacecraft baseline configuration.


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Figure 19. THESEUS Satellite Baseline Configuration and Instrument suite accommodation


The Payload Module is constructed around the IRT instrument, which is partially embedded inside the PLM structure and aligned with respect the S/C symmetry axis. The module top plane is the mounting base of the other instruments DUs (SXI and XGIS) which are distributed around the IRT axis in order to minimize the satellite Moment of Inertia (MoI) but also to support efficient load transfer from the spacecraft to the launch vehicle, respecting their accommodation constraints and thermal requirements.

The IRT LoS is the reference of the overall payload. The telescope is protected from straylight by a baffle; the profile of the baffle defines the position of the entrance window plane of the other instruments.

The 4 SXI DUs are nominally mounted on the opposite side of the solar panels in order to keep them in the coldest side of the satellite and to have the largest area of the observable sky when Theseus lies between Sun and Earth. The SXI DUs positions and orientations are determined in such a way that no X-rays reflected by IRT tube or any other satellite structure can enter into SXI FOV. The 3 XGIS units are tilted in such a way that the FOV of the units partially overlap. The resulting overall FOV (i.e. that of the combination of the 3 units) covers and center the FOV of the 4 SXI modules.


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Figure 20. Structure of THESEUS Payload Module


The proposed accommodation is shown in Figure 22. This configuration guarantees the required nominal SXI, XGIS and IRT Field of view combination (REQ-MIS-050 ). Solar array proposed baseline consists of modular panels in fixed configuration, the panels are composed of a photovoltaic module and a mechanical on which photovoltaic module is mounted. At the level of the payload module top plane a Sun shield is mounted, in order to protect the instruments from solar radiation. The structure of the Payload module is provided of reinforcing shear panels and of an internal cylinder for the IRT telescope and detector accommodation.

The internal cylinder has the function of structural support and it provides also a thermal separation for IRT instrument from the rest of the module. In order to assure radiative thermal decoupling, the telescope is covered by multi-layer insulation.

At the level of focal plane of IRT, where the Detector is placed, a Miniature Pulse Tube Cooler (MPTC) system is provided, acting with a heat-pipe system and a dedicated radiator. Dedicated radiator is provided also for each of the other instruments (SXI and XIGS) on the pedestals. Every radiator is positioned on the own support structure of the single instrument. The Service module is currently conceived as composed by a single module accommodating all the platform units (with the exception of one boresight star tracker which is integrated on the IRT telescope tube to minimize relative misalignment improving pointing knowledge performances. In addition to the platform star tracker, two star trackers hare foreseen in support to IRT instrument, to allow astrometric measurements independent from the system. Figure 21 shows the configuration of Service module while Figure 22 shows Satellite accommodation within the VEGA C fairing.


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Figure 21. THESEUS Thermal concept (Left) and Service Module (Right)


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Figure 22. THESEUS accommodation within VEGA C fairing