The long slit spectrograph onboard the World Space Observatory-- Ultraviolet

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1 The long slit spectrograph onboard the World Space Observatory-- Ultraviolet Maohai Huang* a, Martin A. Barstow b, Zhiyuan Chen a, Jon S. Lapington b, Mikhail E. Sachkov c, Boris Shustov c, Qian Song a, Sen Wang a a National Astronomical Observatories of CAS, 20A Datun Road, Beijing , China b Dept. of Physics & Astronomy, Univ. of Leicester, Univ. Road, Leicester LE1 7RH, UK c Institute of Astronomy of RAS, 48 Pyatnitskaya St., , Moscow, Russia ABSTRACT The World Space Observatory - Ultraviolet (WSO-UV) is a space astronomy project led by Russia, with contributions from China, Germany, Italy, Spain, United Kingdom and a number of other countries in the world. WSO-UV consists of a 1.7-meter diameter telescope and three focal plane science instruments. The Long Slit Spectrograph instrument onboard WSO-UV will produce moderate spectral resolution (R= ) spectra in the 102nm ~ 320nm wavelength range along a slit of 75 arcsec in length and 1 arcsec in width. The spatial resolution of the instrument will be ~1 arcsec. A two-channel scheme is proposed to optimize performance, with each of these using a Rowland Circle optical design with Microchannel Plate detectors in the focal plane. Based on preliminary results of feasibility study we will discuss the detailed design of the spectrograph and its expected performance in this paper. Keywords: Ultraviolet, spectrograph, space telescope, WSO-UV, instrumentation *mhuang@bao.ac.cn; phone ; fax ; BAO.AC.CN 1. INTRODUCTION The World Space Observatory Ultraviolet (WSO-UV)is a 1.7-meter space telescope project [1][2][12] included in the Federal Space Program of Russia for period , and is being developed by an international collaboration effort with contributions from China, Germany, Italy, Spain, United Kingdom and a number of other countries in the world. The telescope is a R-C system with Aluminum+MgF 2 reflective coatings. The science instrument payloads are envisaged to include an imaging instrument in UV and possibly optical wavelength ranges, a high resolution spectrograph with a spectral resolution of working in 115nm ~ 320nm wave length range, and a low spectral resolution spectrograph with one-dimension spatial sampling capability in matching wavelength ranges[3]. This paper describes the low spectral resolution instrument, the Long Slit Spectrograph (LSS) which is undergoing its Phase A feasibility study lead by the National Astronomical Observatories of the Chinese Academy of Sciences. The Sec. 2 of this paper describes main requirements and design constraints to the LSS. Sec. 3 describes preliminary configuration and the optical design. Sec. 4 describes LSS detectors. Sec. 5 gives performance estimation based on current design. 2. MAIN REQUIREMENTS AND CONSTRAINTS 2.1 Main requirements The basic requirement of the LSS is to have a 75 x1 slit spectrograph covering a wavelength range of 115nm to 320nm, with a spectral resolution in the general range of 1000~2500. The baseline requirements during Phase A study are derived from a principal requirement that the Long Slit Spectrograph is to be able to observe faint extra-galactic sources at the far-uv end of the spectral range with high efficiency. The science observing modes of the LSS shall include 1) a time-tagged mode whereas photon events are tagged with a time stamp; 2) an integration mode whereas detector signals are integrated over a given period of time across the spectral range; Space Telescopes and Instrumentation 2008: Ultraviolet to Gamma Ray, edited by Martin J. L. Turner, Kathryn A. Flanagan, Proc. of SPIE Vol. 7011, 70111Y, (2008) X/08/$18 doi: / Proc. of SPIE Vol Y-1

2 3) a high time-resolution (HTR) mode, which performs a series of short, time-stamped integrations. The main user requirements for Phase A are listed in table 1. The in-orbit life time of LSS is 5 years (10 years goal). Table 1. Phase A study baseline user requirements. Main User Requirements Spectral Range 102nm ~ 320nm Spectral Resolution 1500 Slit Length 75 Slit Width 1 ~1.5 Time resolution in HTR mode 1 second Nominal in-orbit life time 5 years 2.2 Major constraints The LSS is to be integrated with the high resolution spectrograph -- the High Resolution Double Echelle Spectrograph (HiRDES), and the fine guidance sensors, to form an integrated monoblock before being integrated with the telescope. The center of field of view on the telescope focal plane is occupied by the imaging unit, the Field Camera Unit (FCU). LSS is to have all its components residing within a boundary corresponding to one-third of the space behind the focal plane. Figure 1 shows the focal plane assembly with the primary mirror unit, the optical bench and its mounting truss, and the focal plane instruments: HiRDES, LSS, and FCU. The focal plane for LSS and HiRDES coincides with the outside boundaries of the instrument facing the optical bench. Any LSS slit also needs to keep clear from the foot-print of the Guiding sensors. OB mounting unit Guiding sensors Optical Bench Long Slit Spectrogra Fig. 1. The focal plane assembly where the LSS resides in WSO-UV. (Image adapted from Lavochkin Association) Besides the usual requirements of mass, power, thermal control, and rigidity, the strict size, shape, and position constraints of the boundaries and slits have been major limiters on the designing of LSS in order to fulfill the user requirements. Proc. of SPIE Vol Y-2

3 3. PRELIMINARY CONFIGURATION AND OPTICAL DESIGN 3.1 LSS Subsystems The LSS has six functional subsystems: the main optics which include the slit assemblies, the gratings and the mirrors with their mountings; the detectors with supporting electronics; the calibrator that provides flatfield and wavelength calibration; the slit-viewer that helps to move the target into the slits; the data processing unit that collects data, controls LSS, and interfaces with the Scientific Data Control Unit on the spacecraft; and the housing and mechanical supporting structures. All subsystems are being studied currently. We describe the optical design and preliminary specifications in the following. 3.2 Optical Design The primary and secondary mirrors of the telescope are Al+MgF 2 coated, so are LSS internal optical elements. Typical reflectivity of a 250Å thick MgF 2 coating over Aluminum are about 80% above 120nm and quickly dropping to about 20% at 102nm[4]. In order to maximize throughput at the far-uv (FUV) end of the required wavelength range, especially that between 102nm and 115nm, only one reflecting surface, which is the dispersion grating, for the FUV wavelength range is used. Due to grating performance variation across the full LSS wavelength range and that different designs of detectors are necessary for FUV and near UV (NUV) photons, a two-channel design is selected for the LSS. The two channels cannot observe the same area of the sky simultaneously. In order to ensure reliability of the LSS over its 5-year nominal life time, no mechanism is needed to operate within the LSS in order to switch observing channel. Both channels of the LSS has its own slit, optics, and detectors. A small repointing is used to direct light of an object from one channel to the other. Figure 2 shows the exposed opto-mechanical layout of the LSS. The NUV channel, the calibrator, and the slit-viewer in the figure are drawn based on preliminary design options. slit-viewer Fig. 2. LSS optical system as well as the detectors, calibrators, and slit-viewer are shown within the allocated space. The NUV channel, the calibrator, and the slit-viewer are drawn based on preliminary designs. Proc. of SPIE Vol Y-3

4 Both the FUV and the NUV channels are designed based on the Rowland Circle with modifications to accommodate requirements and constraints. The FUV Channel uses a spherical concave holographic grating as the dispersion element, which is the only reflecting surface between the slit and the FUV detector. In order to move the detector into the boundary box of the LSS the distance between the grating and the detector (837mm) is shorter than that between the slit and the grating (1000mm). The NUV Channel has 3 to 4 reflective surfaces including the concave grating between the slit and the NUV detector. Because the reflectivity of the MgF 2 coating is much higher in the NUV range than in FUV a number of pick-up and redirecting mirrors can be used in order to place both slits and all LSS internal components within the geometrical constraints without degrading the throughput seriously. 4. DETECTORS The FUV and the NUV channels of LSS have their own Microchannel Plate (MCP)detectors. The type of detector is mainly determined by the working wavelengths. We choose MCP detectors with a Caesium Telluride (Cs 2 Te) photocathode for the NUV channel and that with a CsI photocathode for the FUV channel to optimize detector efficiency. The 102 nm lower wavelength cutoff of the FUV detector necessitates using open-faced design, and quantum efficiency in the wavelength range will be maximized by use of a Caesium Iodide (CsI) photocathode deposited on the input surface of the microchannel plate. Since CsI can tolerate only limited exposure to atmosphere, the detector has a vacuum door which is opened for ground testing, manually reclosed and finally deployed for in-orbit operation. The wavelength range of the NUV detector allows a sealed tube design with input window to be used, removing the complexity and increased risk associated with a vacuum door, and simplifying operation of the detector both on ground and in orbit. A semi-transparent Caesium Telluride photocathode is deposited on the inside of a window made of Magnesium Fluoride to accommodate the wavelength range. The LSS optical design has engineered to match with commercially / available detector body and MCP dimensions, avoiding a highly custom design and reducing costs substantially. The high aspect ratio imaging area required by the spectrometer will be accommodated within a circular detector with a nominal 75 mm active diameter. The detector will use a small pore MCP stack operating at a gain of ~10 7 electrons. The detector will be read out using the Image Charge technique[5], a method whereby the event charge is collected on a resistive layer deposited on a dielectric window, which serves to instantaneously localize the event whilst allow it to discharge over a longer timescale. The transient arrival of the event pulse is detected by a centroiding charge division readout, such as the Vernier anode[6], which is capacitively coupled to the back side of the dielectric window. This technique enhances performance by mitigating distortions produced by redistribution of secondary electrons produced on the anode, and greatly simplifies the tube manufacture by allowing the readout to be outside the tube vacuum envelope. Several designs of centroiding readout schemes are being considered, the Vernier anode being the current baseline option. Magnesium Fluoride input window Cs2Te photocathode "Image Charge" resistive layer MCP stack N Dielectric window Charge division image readout I Detector vacuum housing Fig. 3. A schematic cross-section of the proposed NUV sealed tube detector showing the individual components. Proc. of SPIE Vol Y-4

5 Table 2. LSS detector parameters. LSS detector parameters NUV channel Wavelength Range 160nm~320nm Total gain ~10 7 Spatial resolution 20 µm Open area ratio (OAR) 60% Active area 60mm x 6mm Read out Method Centroiding charge division Photocathode material Cs 2 Te Quantum efficiency 20% FUV channel Wavelength Range 102nm~170nm Total gain ~10 7 Spatial resolution 40 µm Open area ratio (OAR) 60% Active area 60mm x 6mm Read out Method Centroiding charge division Photocathode material CsI Quantum efficiency 25% The background noise from the FUV detector should be less than 0.01 count/cm 2 /s using low noise MCPs, and for the NUV detector the Cs 2 Te gives higher dark current[4]. The combined background from real in-orbit measurement tends to be much higher. 5. ESTIMATED PERFORMANCE 5.1 Parameters and assumptions Although some design choices that could have impact on NUV channel efficiency haven t been made, based on current optical designs and detector performances we are able to estimate the brightest and the faintest sources the LSS could be expected to observe. Table 3 summarizes available or assumed LSS parameters for its current design and the estimated minimum flux and maximum flux that can be observed. The primary mirror area of the WSO-UV T-170M telescope is 2.27m 2. Due to blockage of the secondary mirror unit the collecting area is lost by 10%. We adopt a reflectivity of 0.8 of all mirrors above 115nm and 0.2 at 102nm[4]. A 0.95 slit efficiency is assigned to both channels. Multiple internal mirrors (3 in the NUV channel) produce the combined efficiency. The grating and detector efficiencies, as well as detector maximum count rate, are taken from early estimations from component providers[7]. Compared with the Cosmic Origins Spectrograph (COS) instrument for the HST[8][9], the sensitivities, which is defined as event count rate generated from a resolution element ( resel hereafter) given unit incident flux, of the LSS channels are comparable given current uncertainties. The effective area at the short wavelength end of the FUV channel is severely affected by the MgF 2 coating but is still better than that of any channel of FUSE s[11]. LSS detector noise figures are calculated based on on-orbit measurement of the background noise flux of FUSE (0.8 count/s/cm 2 [10]). Compared with FUSE and HST, WSO-UV will operate at a much higher orbit (35860 km GSCO) in a Proc. of SPIE Vol Y-5

6 thinner exosphere with considerably less earth shine problem and possibly much lower geocoronal airglow interference, but much less night observing time. The sky background of LSS will be further analyzed in later works. We ignore the sky background in this paper. Table 3 LSS parameters and performance estimation LSS Parameters and Performance Estimation NUV (160nm~320nm) Combined efficiency of internal mirrors 0.51 Grating efficiency 0.3 Detector efficiency 0.2 Maximum detector count rate 10 5 s -1 Effective area 378cm 2 Sensitivity 7x10 13 ct/resel/s/(erg/cm 2 /s/å) Detector background flux 5x10-19 erg/cm 2 /s/å/resel 10 hr SNR=10 flux average 4x10-17 erg/cm 2 /s/å Maximum observable flux 1x10-12 erg/cm 2 /s/å Slit dimension 1 x 75 Spectral resolution 1.5 Å(TBC) Spectral resolving power at channel center 1600(TBC) Spatial resolution ~1 arcsec(tbc) FUV long wavelength end (115nm~170nm) Combined efficiency of internal mirrors 1 Grating efficiency 0.24 Detector efficiency 0.25 Maximum detector count rate 10 5 s -1 Effective area 745cm 2 Sensitivity 4x10 13 ct/resel/s/(erg/cm 2 /s/å) Detector background flux 1x10-18 erg/cm 2 /s/å/resel 10 hr SNR=10 flux average 7x10-17 erg/cm 2 /s/å Maximum observable flux 3x10-12 erg/cm 2 /s/å Slit dimension 1 x 75 Spectral resolution 0.8 Å Spectral resolving power at channel center 1700 Spatial resolution 1 arcsec FUV short wavelength end (102nm~115nm) Combined efficiency of internal mirrors 1 Grating efficiency 0.14 Detector efficiency 0.25 Effective area 27cm 2 Sensitivity 1x10 12 ct/resel/s/(erg/cm 2 /s/å) Detector background flux 4x10-17 erg/cm 2 /s/å 10 hr SNR=10 flux average 2x10-15 erg/cm 2 /s/å Maximum observable flux 1x10-10 erg/cm 2 /s/å 5.2 Faintest and brightest sources observable The 10 hr SNR=10 flux average in Table 3 shows the average flux that can be observed at SNR=10 after a 10 hour integration, assuming the detector Poisson noise discussed above is the only source to limit detection of faint sources (i.e. Proc. of SPIE Vol Y-6

7 in the detector noise dominated regime, rather than the source photon dominated regime). These fluxes in Table 3 means that a 23m V, 24 m V, and 21 m V B0 star can be observed by the NUV channel, the long end of the FUV channel, and the short end of the FUV channel, respectively, at SNR=10, in 10 hours. In the detector noise dominated regime, the minimal observable flux for LSS NUV channel are better than that of COS G230L channel because LSS NUV detector is not expected to suffer from the impurity problem in MAMA detector window[8]. Detector noise dominated detection limit of LSS FUV channel is expected to be about the same as the COS FUV channels, and is better than that of FUSE[11] at the short wavelength end of FUV channel. At 102nm the resolution element area on the LSS detector is 40µm x 40µm for LSS, much smaller than that of FUSE LiF1 channel re-sampled to such resolution (~900µm x 50µm). Because the intrinsic background of an MCP detector is proportional to its active area, and we assume LSS detectors to have similar noise event rate per unit area compared with COS and FUSE, the smaller areas of LSS resolution element on its detectors generate less noise, and results higher capability to detect faint sources. It needs to be pointed out that, like most of its peers, LSS on-board WSO-UV is photon-starved for observing faint sources, especially point sources. For example if one is to observe the faintest point source that matches the detector noise floor (10-19 erg/cm2/s/å/resel level) of LSS at SNR=3, thousands of hours of integration time will be needed to collect enough photons. However this problem is alleviated for extended sources, because the long slit of LSS allows improvement of LSS throughput at the cost of spatial resolution, through re-binning resolution elements in the cross dispersion (i.e. spatial) direction. For example a ten-resel (~10 arcseconds) re-binning would reduce the integration time by an order of magnitude. Therefore with its slit design LSS has a greater advantage in observing very faint extended objects. The brightest source that can be observed by LSS is limited by the maximum count rate of the photon-counting MCP read-out unit. We assume 10 5 count/sec in our calculation. The brightest B0 stars LSS can observe are then ~12m v for the full range except for the short wavelength end of FUV, for which they are ~9 m v. 6. CONCLUSIONS We present the preliminary Phase A study results of the Long Slit Spectrograph to for the World Space Observatory Ultraviolet mission. The LSS uses modified Rowland Circle design in its two observing channels for the NUV (160nm ~ 320nm) and the FUV (102nm ~ 170nm) wavelength ranges. The design centers around maximizing sensitivity at the short wavelength end of the FUV channel. CsI and Cs 2 Te photocathode Microchannel Plate detectors are used for the FUV and NUV channels, respectively. There is no moving mechanism needed for switching between the two channels during a full spectral range observation. The slit size is 75 x1 and the average spectral resolving power is ~1500. Based on current effective area and detector noise data, LSS is capable of observing flux level better than erg/cm 2 /s/å per resolution element (SNR=10, 10 hour integration) for all but the shortest part of spectral ranges. LSS is expected to have similar sensitivity compared with COS low dispersion channels, and should perform better than COS G230L channel toward the detector noise dominated regime. LSS could observe between 102nm and 110nm which is inaccessible to COS. At the short wavelength end of FUV range LSS sensitivity and noise level are better than FUSE (single channel re-binned to matching spectral resolution). LSS has advantages in observing faint extended objects due to its slit compared with COS. It should be noted that scattered light, space environment effects, and other system uncertainties all contribute to the background and calibration of a real observation and have impact on instrument performances. Some of the negative effects could be difficult or impossible to remove. Our estimations above are based on our current understanding of our system with many uncertainties ignored. The results here should be taken as reference points rather than definitive specifications. ACKOWLEDGEMENTS: We thank Fu-Zhen Cheng, Xiaowei Liu, Suijian Xue, Tinggui Wang, Lihong Geng, Caihong Sun, Jinxin Hao, and Gang Zhao of the WSO-UV Chinese working group, and the WSO-UV international team, especially Norbert Kappelmann, Klaus Werner, Isabella Pagano, Anna Ines Gömez decastro, Alexander Moisheev, and Willem Wamsteker, for their support and input to LSS study. We are grateful to Zhang Wei of JHU for helpful Proc. of SPIE Vol Y-7

8 discussions on LSS performance comparison with COS. LSS study is partly supported by the National Science Foundation of China grant number REFERENCES [1] [2] [3] [4] [5] [6] [7] [8] [9] [10] [11] [12] Shustov, B., Sachkov, M.; Gömez decastro, A.I., Huang, M., Werner, K., Kappelmann, N., Pagano, I., "WSO-UV - Ultraviolet Mission for the Next Decade," Astrophysics and Space Science, in press (2008). Yaskovich, A., Zverev, A., and Shustov, B., "Optimization of the system consisting of the telescope and the focaldevice unit in the Spectrum-UV Project," J. Opt. Technol., 73, (2006). Pagano, I., et al, "The focal plane instruments on board WSO-UV," Astrophysics and Space Science Proceedings series, M. Chavez, E. Bertone, D. Rosa-Gonzalez & L. H. Rodriguez-Merino (eds.), (2007). Madden, R.P., [Phys. Thin Films], Academic Press Inc., New York, Vol. I, , (1963). Jagutzki, O., Lapington, J.S., Worth, L.B.C., Mergel, V., and Schmidt-Böcking, H., Position sensitive anodes for MCP read-out using induced charge measurement, Nucl. Instr. Meth., A477, 256 (2000). Lapington, J.S., Sanderson, B., Worth, L.B.C., and Tandy, J.A., Imaging achievements with the Vernier readout, Nucl. Instr. and Meth., A477, 250 (2002). Lapington, J.(private communication). Soderblom, D. R. et al., [Cosmic Origins Spectrograph Instrument Handbook], STScI, Baltimore, version 1.0, (2007). Green, J.C., Morse, J.A., Andrews, J.P., Wilkinson, E., Siegmund, O.H.W., Ebbets, D. "Performance of the Cosmic Origins Spectrograph for the Hubble Space Telescope," Ultraviolet-Optical Space Astronomy Beyond HST conference proceedings, eds J. A. Morse, J. M. Shull, A. L. Kinney, Astronomical Society of the Pacific Conference Series, 164, p (1999). Sahnow, D. J., et al, "ON-ORBIT PERFORMANCE OF THE FAR ULTRAVIOLET SPECTROSCOPIC EXPLORER SATELLITE," The Astrophysical Journal, 538, L7 L11 (2000) Andersson, B-G. (ed), The FUSE Observer's Guide,, V8.0, Sachkov, M., World Space Observatory-Ultraviolet: International Mission for UV Spectroscopy and Imaging, in VI Serbian Conference on Spectral Line Shapes in Astrophysics (VI SCSLSA), AIP Conference Proceedings, Vol. 938, (2007) Proc. of SPIE Vol Y-8