The T170-M telescope

The heritage for the WSO telescope design 1 is the Russian-led international space observatory ASTRON, launched in 1983. It functioned for 6 years and was the first UV telescope placed into a highly eccentric orbit. This led to the proposed development of the Spectrum UV mission with Russian, Ukrainian, German and Italian partnership. However, the major social and economic changes taking place in the former Soviet Union countries has reduced funding of the Russian space programme and limited progress in all space projects. The experience gained in the development of the Spectrum UV T170 telescope is, therefore, now being applied to a modified version (the T170M) for the much broader international scope of the WSO.

The structure of the T170M telescope is shown in Fig. 1, with the principle structural elements; the Primary Mirror Unit (PMU), the Secondary Mirror Unit (SMU) and the instrument compartment. The SMU is attached to the telescope structure with a spider. The optical design is a Ritchey-Chretien type with a 1.7m hyperbolic primary mirror and engineering model of this unit for the T170 design (with a spherical Sitall mirror) has successfully passed through a number of environmental tests. Modifications for the WSO centre on reducing the overall mass of the system below 1650kg. Consequences of this are a shorter relative focal length, although the f number (10) remains the same, and reduced distances between the mirrors and between the primary and optimum focal surface (Table. 1).




Aperture diameter



Telescope focal length



R 1 (primary radius)



e 1 (eccentricity)



SM light passing diameter



R 2 (secondary radius)



e 2 (eccentricity)



Space between mirrors



Obscuration with optimum baffles



Space between PM apex and optimum focal surface



Table 1. Summary of the technical details of the T170M telescope design and comparison with the earlier T170.




 Fig. 1: Schematic diagram of the WSO 1.7-m aperture UV telescope.


The secondary mirror is design includes capabilities for in-flight alignment and focusing. These are controlled by means of: 1) a 1mm displacement of the secondary along two mutually perpendicular axes (in a plane normal to the optical axis) in steps of 0.01mm; 2) focusing with axial movements of the secondary mirror within ±5mm from the normal position, in steps of 0.002mm. Both primary and secondary mirrors will be coated with Al and a MgF 2 layer, to yield an interference maximum at 115nm.

The primary mirror is supported by a central mount and includes a system of heaters on its lower side to minimize the thermal gradients along the mirror during operations. A support structure consisting of 12 profiles Invar bars and is covered by corrugated Al sheets to reduce thermal distortions of the tube. An rms wavefront error of less than 13nm is required of the system to deliver diffraction limited optical performance. Figs 2. and 3. show the predicted telescope performance as a function of distance from the centre of the field. The diffraction limit corresponds to 0.035 arcsec.



The High Resolution Double Echelle Spectrograph (HIRDES)

The UV spectrometer 2 (Fig. 4) comprises three different single spectrometers. Two of these are echelle instruments, designed to deliver high spectral resolution, and the third is a low dispersion long slit instrument (LSS). At high dispersion, the 110 to 320nm waveband of the WSO will be divided into two, the UV (UVES, 178-320nm) and VUV (VUVES, 103-180nm). The fundamental concept of HIRDES is based on the design heritage of the ORFEUS missions (Orbiting and Retrievable Far and Extreme Ultraviolet Spectrometer, mounted on the ASTRO-SPAS free flyer), successfully flown on two STS flights in 1993 and 1996.

Each of the three spectrometers has its own entrance slit, lying in the focal plane of the T-170M telescope on a circle with diameter 100 mm. The three optical trains are not used for simultaneous observation, but in sequential mode. This is managed by satellite motion with a pointing stability requirement of 0.1 arcsec to be monitored by three Fine Guidance Sensors. Each of the three sub-instruments includes optical elements to form the spectral imaging and uses a main and a redundant detector (baseline: MCP/WSA detector) for observation. The main and the redundant detectors are placed closely together in a L-shaped detector housing. A servo-driven mirror in front of the detectors will illuminate either of the detector apertures. Between this servo mirror and the redundant detector an additional servo-driven gray filter will be fitted, to reduce the intensity, so the redundant detector can be used to observe brighter objects if the front-end electronics of the main detector is saturated. Technical details of the spectrographs are listed in tables 2, 3 and 4, while the optical layout is in Fig. 4.


Entrance aperture

circular 80 m m

Collimator mirror

toroidal, R1 = 1608 mm, R2 = 1593 mm, circular 80 mm

Echelle grating

40 grooves/mm, 66.9° blaze, 90 mm x 190 mm

Cross dispersion prism

12° quartz, double pass, 100 mm x 110 mm

Camera mirror

spherical, R=1600 mm, 110 mm x 130 mm


30 mm (echelle dispersion), 30 m m resolution (pixel-width)

40 mm (prism dispersion), 40 m m resolution (pixel-height)

Resolving power

R = 50000 (3 pixel criterion)

Table. 2: Technical details of the UV echelle spectrograph (UVES)

Entrance aperture

circular 80 m m

Collimator mirror

parabolic, R = 1600 mm, 6°-off-axis, circular 80 mm

Echelle grating

65 grooves/mm, 71° blaze, 90 mm x 270 mm

Wadsworth grating

toroidal, 625 grooves/mm, R1 = 1698 mm, R2 = 1695 mm, 90 mm x 120 mm


30 mm (echelle dispersion), 30 m m resolution (pixel-width)

40 mm (Wadsworth dispersion), 40 m m resolution (pixel-height)

Resolving power

R = 55000 (3 pixel criterion)

Table. 3: Technical details of the VUV echelle spectrograph (VUVES)

Entrance aperture

rectangular, 80 m m x 6 mm

Rowland grating

aberration corrected concave grating, R » 800 mm


80 mm (grating dispersion), 15 m m resolution (pixel-width)

Resolving power

examples: R = 1000 (3 pixel criterion), 110 nm to 350 nm

Table. 4: Technical details of the long slit spectrograph (LSS)


The original spectrometer design for ORFEUS incorporated a Z-stack MCP detector with a wedge-and-strip anode readout. The leading two MCPs are standard resistivity plates with 12.5�m pores, 15 o bias angles (for HIRDES) and L/D ratio of 80:1, while the rear MCP is a low resistance L/D 40:1 unit. The useful detector area is 40mm x 30mm. The spatial position of the impinging photons corresponds to the focal point of the generated electric charges at the third MCP outlet and is detected by means of the wedge-and-strip anode. From four electrodes the charges of an electron cloud are measured with ADCs with an oversampling technique. From these four charge values the exact coordinates of the centre of the electron cloud are calculated in the ICU. The spatial resolution in the WSA plane depends on the ADC resolution. With 12 bit ADCs a minimum resolution of 30 m m (x-direction) and 40 m m (y-direction) can be achieved, corresponding to 1024 x 1024 pixels in the detection plane. With 14 bit ADCs a minimum resolution of 7 m m (x-direction) and 10 m m (y-direction) can be achieved, corresponding to 4096 x 4096 pixels. These resolutions are well matched to the spectrometer design requirements.

Adoption of a common detector design for the long slit dispersion spectrometer poses a surprising problem. Although the instrument operates at lower spectral resolution, the full 110-350nm spectral range cannot be accommodated on a single detector of the above dimensions. Hence, there is a potential trade-off between the desired resolving power (R~1000) and spectral coverage. One possible solution is to use the higher spatial resolution detector design, to allow the full wavelength range to be compressed into a smaller detector field. We propose to adopt an MCP detector design, using a Vernier anode readout 3 , which has recently been proven in-flight on board the J-PEX high resolution EUV spectrometer 4,5 .



WSO performance with the HIRDES spectrographs

To be a viable replacement for HST, the overall capabilities of the WSO prime spectroscopic instruments need to compete with the performance of the STIS and COS spectrographs, which they will eventually supercede. In having a mission dedicated to a narrower waveband than HST, which also has coverage extending into the visible and infrared, there are several advantages. First, a simpler optical layout of the spectrograph systems is available. A resulting reduction in the number of necessary reflections, to accommodate the instruments in the available space, leads to lower light losses. Second, all the available observing time can be devoted to UV astronomy. Third, the intention of placing the telescope at the L2 Sun-Earth point yields a more efficient duty cycle, compared to low Earth orbit. Fig. 5 shows the comparison in the effective area of HST with the STIS spectrograph and the WSO UVES and VUVES instruments, while Table 5 compares the WSO performance with that COS.

WSO has an effective area ~4-10 times that of HST/STIS, while providing a resolving power approaching that of the highest STIS echelle capability. When this gain is folded with the improvements in available observing time, the overall productivity will be a factor 40-50 higher than STIS. In contrast, at the shorter wavelengths the effective area of COS is larger than VUVES (although again the observing time factors and the larger wavelength coverage make the productivity comparable) but falls below the performance of UVES at the longer wavelengths. However, the COS resolving power is a factor of two lower.


? (nm)


A eff (cm 2 )

? (nm)


A eff (cm 2 )






















Table. 5: Comparison of the predicted performance of the WSO UVES and VUVES spectrographs (right) with the HST COS instrument (left).


WSO focal plane imagers

Although the primary science of the WSO mission is spectroscopic, there is an important role for high spatial resolution UV imaging of the sky. Therefore, it is planned to include a complement of UV imaging detectors 6 in the focal plane, to provide serendipitous science during spectroscopic observations as well as planned studies of specific target areas. Much of the available volume in the focal plane, immediately behind the primary mirror, is occupied by the HIRDES. This leaves only a very narrow space (10cm diameter cylinder) on the telescope axis that can be used for a direct imager, which samples the best diffraction limited resolution of the optical system. This detector (the central field camera, CFC) will require a fixed filter, since there is no room for a filter wheel.



Fig. 6: Side view of the optical bench, showing three UV imagers, the redirecting prism and one possible arrangement for the fine guidance sensor.


Fig. 7: Proposed focal plane optical bench, accommodating 6 UV imagers in addition to the central field camera. The spectrograph slit positions are also shown




0.03 (2.5�m)



0.15 (12.5�m)





Table. 6: Performance of UV and optical imagers.

One solution to the accommodation problem for additional imagers is to redirect the beam through 90 o by a single reflection, as shown in figs. 6 and 7. An option currently under investigation is to operate 3 imagers at f/10, each with a flat mirror pickup, and 3 at f/50, using magnifying relay of two mirrors. Thus the f/50 cameras would operate at optimal spatial sampling and there would be more space in the aperture to accommodate the fine guidance sensors. These HSI cameras are based on the TAUVEX design for Spectrum X-gamma and each is equipped with a 4-element filter wheel. Combining two wheels with 3 filters plus one open position will provide 6 colours plus a completely open position. As can be seen from table 6, which summarises the target performance for the imagers (for a plate scale of 12.06 arcsec mm -1 ), the demands on detector spatial resolution are stringent. The J-PEX detector design, discussed above for the LSS, provides the best available but is ultimately limited by the MCP pore size of 6�m.

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