2.2.1.

MISC-T operation principle

The MISC-T module is primarily dedicated to pointed observations to detect bio-signatures in habitable worlds in primary and secondary exoplanet transits. MISC-T has three channels (T-S, T-M, and T-L) that share the same field-of-view (FOV) by means of beam splitters, and all channels are operated simultaneously to cover the full spectral range from 2.8 to $20\mu \mathrm{m}$ at once. The science traceability matrix requirements for the spectral accuracy for exoplanet transit spectroscopy are $<5\mathrm{ppm}$ from 2.8 to $10\mu \mathrm{m}$ and $<20\mathrm{ppm}$ from 10 to $20\mu \mathrm{m}$. To reach this performance requires system optimization involving not only the MISC-T optical design and detector/readout electronics choices, but also observatory performance issues such as the pointing performance and the nature and variability of the Origins telescope aberrations. Various instrument design concepts were considered during the present study. The instrument design will be finalized in Origins Phase A. For proof-of-concept and costing purposes, we designed MISC-T with a novel densified pupil optical design, and HgCdTe and Si:As detector arrays. The densified pupil design divides up the light from a point-source in the instrument pupil plane up into a number of individual spectrometer optical channels and is described in more detail in Ref. 5. This design is relatively insensitive to minor line-of-sight pointing drifts and telescope aberrations, and the detectors do not require a sub-Kelvin refrigerator. It does have the disadvantage of reduced sensitivity to faint sources compared to the traditional spectrometer design since the point source spectra are spread out over a relatively large number of detector pixels. However, for this science program with the relatively bright host star targets, spectro-photometric stability is more important than faint-source sensitivity.

To demonstrate the potential of this novel optical design, a testbed system has been constructed using the densified pupil design and a $1024\times 1024$ Si:As Impurity Band Conduction array. Performance testing has been ongoing at NASA Ames Research Center and so far, the design has met all of its design requirements.

It is noteworthy that an alternative instrument design concept was considered in which transition edge sensor (TES) detectors were fed by a more conventional spectrometer, with Winston cones used to concentrate more photons to improve sensitivity at these short wavelengths. As this design concept has not been analyzed as thoroughly for overall system performance, especially with respect to errors introduced by telescope image quality and pointing variation, it was not chosen as the baseline, but further work on this alternate concept will be carried into the Phase A Concept Study.

Densified pupil spectroscopy is a new optical system approach for transit spectroscopy^5,6 that will greatly improve spectro-photometric performance in the presence of optical disturbances. With pupil densification, the MISC-T science image will not be disturbed by modest telescope pointing jitter (e.g., 10 mas) or telescope aberrations. The large number of science pixels minimizes the effects of intra- and inter-pixel sensitivity variations. Reference pixels also provide potential detector baseline common-mode gain fluctuation calibration using a cold photon shield mask. As part of the Origins technology development plan, the team will refine an optimal pixel-to-pixel calibration technique using an ultra-stable internal calibration source (see Sec. 2.2.3.) that meets with the instrument requirement of $<20\mathrm{ppm}$ at $\lambda =10.5$ to $20\mu \mathrm{m}$. The MISC-T block diagram is shown in Fig. 1.

Fig. 1

The block diagram of the MISC-T.

No simultaneous operation between MISC-T and other instruments is expected; therefore, the high operating temperature for the HgCdTe detectors will not be a source of noise to the other instruments.

A Lyot-coronagraph-based tip–tilt sensor (TTS) located in the instrument fore-optics uses the light reflected by a field stop, which corresponds to 0.3% of the light from the target, to send fine pointing information to the fine steering mirror (FSM) in the telescope optical train.⁷ The TTS measures the long-term drift using the point source image formed on the focal plane through a Lyot stop and feeds this information back to the observatory to mitigate the long-term drift down to 2 mas over a timescale of a few hours.

2.2.2.

MISC-T optical design

The MISC-T optical design is based on the densified pupil spectroscopy.⁵ Conventional spectrographs (e.g., long slit) use focal plane division: light from multiple positions in the focal plane is dispersed into multiple spectra. In contrast, the densified pupil spectrograph uses pupil-plane division and produces spectra from multiple sub-apertures of the pupil plane. It should be emphasized that in spite of the novel arrangement of optics in the densified pupil design, the optical components themselves are completely ordinary. The performance of this design is currently being demonstrated in tests at NASA Ames Research Center. The densified pupil spectrograph is composed of a pupil division/densification system followed by a normal spectrograph system (See Fig. 2).

Fig. 2

The MISC-T Optical Design. The three spectrometers are fed by a collimated beam that emerges from the pupil slicer assembly and the TMA collimator. The dimensions of the instrument in mm are indicated.

The pupil densification plays a key role in performing highly stable spectroscopy. Because of the pupil division into five slices in the case of MISC-T, and the densification of the divided pupil, each of the densified sub-pupils acts as a point source, and each corresponding beam is collimated by triple mirror assembly (TMA) collimator as shown in Fig. 2. The pupil is sliced into five segments, as shown in Fig. 3.

Fig. 3

The telescope pupil is reimaged and sliced as shown at left, with the slices separated as shown in the image on the right.

The first dichroic beam splitter reflects the light of 11 to $20\mu \mathrm{m}$, which is led to the spectrograph of T-L, and transmits the light of $<11\mu \mathrm{m}$. The second dichroic beam splitter reflects the light of 5.5 to $11\mu \mathrm{m}$, which is led to the spectrograph of T-M, and transmits the light of $<5.5\mu \mathrm{m}$, which is led to the spectrograph of T-S. Field stops with radii of 3.″0, 3. ″0, and 2. ″0 are installed in the spectrograph T-S, T-M, and T-L channels, respectively, to reduce the zodiacal background photons. The field stop aperture size determines the FOV size, and no spatially resolved information within the FOV is obtained.

Each spectrograph employs a transmission grism to disperse the light in one dimension as a function of the wavelength on a detector plane. Each of the five spectra shares 271 pixels in the dispersion direction by 27 pixels in the cross-dispersion direction for the T-S, 551 pixels in the dispersion direction by 54 pixels in the cross-dispersion direction for the T-M, and 1586 pixels in the dispersion direction by 65 pixels in the cross-dispersion direction for the T-L. The layout of these spectra onto the detector arrays is shown in Fig. 4. The linear dispersion ($\mathrm{\Delta }\lambda$) is set to 0.054, 0.108, and $0.074\mu \mathrm{m}$ for T-S, T-M, and T-L, respectively.

Fig. 4

The spectra from the five pupil slices are overlaid on the (a) short; (b) mid; and (c) long focal planes.

The TTS corrects any long-term drift in the telescope pointing. We conducted a trade study between a Lyot-Coronagraph-based field stop with a dedicated TTS array versus. a guider camera with a dichroic beam splitter to extract 2 to $2.8\mu \mathrm{m}$ light of the target. The study showed a clear advantage in the Lyot-Coronagraph method in terms of MISC-T optical throughput, and in simplicity and cost. The TTS uses only 0.3% of the light from the target host star, which is reflected by the focal plane mask in the fore-optics, and the majority ($\sim 98\%$) of the light from the target passes through $\sim 5^{″}$ aperture in the focal plane mask and is led to the T-S/M/L spectrometer channels. The long-term drift is measured by examining the position of the point-spread function (PSF) centroid of the image produced on the detector plane of the TTS through the Lyot stop. Assuming that the actual Lyot stop size is the same size as the Origins pupil and has an accuracy of 1%, the TTS will work for targets of M-type stars brighter than $K=10\mathrm{mag}$. The TTS feeds back this information to the observatory, and the long-term drift can be mitigated down to 2 mas for M-type star brighter than $K=10\mathrm{mag}$. Figure 5 shows the MISC-T TTS schematic.

Fig. 5

A schematic view of the MISC-T TTS showing the results of a simulation of a PSF image (a) at the Lyot-Coronagraph focal plane mask; (b) just before the Lyot Stop; (c) just after the Lyot stop; and (d) at the detector plane.

2.2.3.

MISC-T detectors and read-out electronics

The MISC-T’s baseline employs two types of detectors: (1) HgCdTe that is bonded to a Si readout multiplexer, and (2) Si Blocked Impurity Band (BIB) design that is bonded to a Si readout multiplexer.⁸ Each of the MISC-T-S and T-M optical paths contains a $2\mathrm{k}\times 2\mathrm{k}$ HgCdTe detector array operated at 30 K, and the MISC-T TTS employs a $256\times 256$ HgCdTe detector array operated at $\sim 30\mathrm{K}$. HgCdTe detectors developed for the Near Earth Object Survey Mission (NEOSM) have successfully demonstrated performance out to $10\mu \mathrm{m}$.⁹ The MISC-T-L optical path contains a $2\mathrm{k}\times 2\mathrm{k}$ Si:As detector array bonded to a Si readout multiplexer operating at $\sim 8\mathrm{K}$ to provide good detective quantum efficiency over the 10.5 to $20\mu \mathrm{m}$ wavelength range. Ultra-stable internal calibration sources are employed to help achieve the required detector stability of $<20\mathrm{ppm}$. These sources are composed of a combination of an infrared source black-body combined with a shorter-wavelength photocell detector used for temperature control. The photocell measures the source output in its blackbody Wien limit, allowing very sensitive feedback of the calibration source temperature, while the infrared wavelengths used for the infrared detector calibration are in the blackbody Rayleigh-Jeans limit. The black-body, whose light path comes from a region just outside the optical path coming from the telescope, is defocused and provides a flood illumination of the array.

Amplification of the signals from the detectors is a straightforward reuse of similar technology from previous space applications of HgCdTe and Si:As BIB detectors, including Spitzer, James Webb Space Telescope (JWST), Wide-field Infrared Space Explorer (WISE) and NEOSM. This method employs dedicated satellite chips located in close proximity to the detectors operating at cryogenic temperatures. The cable run from the cold satellite chips near the detectors to the warm electronics in the Origins design is shorter than it is in JWST, so the difficulties that were experienced with the long cable runs in the latter mission should be much reduced. MISC-T incorporates flexible readout patterns and a variety of readout strategies including double-correlated and Fowler sampling techniques. Due to the extreme detector stability requirements, the readout electronics will need to be carefully designed and stabilized in temperature.

2.2.4.

MISC-T observing mode and data rates

MISC-T has a single observing mode: AOT01 (Table 2). MISC-T-S, T-M, and T-L share the same FOV by means of beam splitters and are simultaneously operated to produce a complete 2.8 to $20\mu \mathrm{m}$ spectrum at once.

Table 2

The MISC observing modes (baseline) and data rates.

2.2.5.

MISC-T thermal and mechanical design

MISC-T instrument components are located in two thermal zones. The instrument optics and detectors are located in a cold zone at the same $\sim 4.5\mathrm{K}$ temperature as the telescope. The detectors have a weak thermal link to this cold temperature so they can be heated slightly and held at a stable operating temperature. This operating temperature is $\sim 30\mathrm{K}$ for the HgCdTe detectors and $\sim 7\mathrm{K}$ for the Si:As detector. Baffling in the cold instrument prevents IR light leakage from the warmer detectors from reaching the SI:As detector. This was successfully demonstrated by WISE, which had 32 K detectors and 8 K detectors located in the same instrument. The drive electronics that handle instrument data collection and communication with the spacecraft are located in the $\sim 300\mathrm{K}$ warm zone.

The MISC-T 3D solid model is shown in Fig. 6. To reduce MISC-T instrument mass, beryllium was assumed as a baseline material for the mirrors and mirror support structures, and the base plate. Cordierite CO-720, a ceramic material by Kyocera, has also been studied as an alternative material for the MISC-T and was used for the test instrument to demonstrate the performance of the densified pupil spectroscope at NASA Ames Research Center. Table 3 lists the mass breakdown of cold component estimated each for beryllium and Cordierite models. MISC-T does not have any moving parts or mechanisms, which is very important to achieve high spectro-photometric stability. Table 4 lists MISC-T resource requirements (volume, mass, and power).

Fig. 6

A 3D solid model of the MISC-T instrument indicating the optical path of the light within the instrument and the location of the mirrors, optical elements and supporting structures on the optical baseplate.

Table 3

Mass breakdown of MISC-T cold component.

Table 4

MISC-T instrument resource requirements.

2.2.6.

MISC-T instrument control

The MISC-T instrument includes redundant dual-string warm electronics boxes (WEB A and WEB B for MISC-T). Each electronics box contains a single board computer (SBC) and three boards: (1) a CPI/ Power board, (2) a digital and analog board for array control (two $2\times 2\mathrm{k}$ HgCdTe arrays for T-S and T-M, one $2\mathrm{k}\times 2\mathrm{k}$ Si:As array for T-L, one $256\mathrm{k}\times 256$ HgCdTe array for T TTS), and (3) a board for peripheral control for detectors (two $2\times 2\mathrm{k}$ HgCdTe arrays for T-S and T-M, one $2\mathrm{k}\times 2\mathrm{k}$ Si:As array for T-L, one $256\times 256$ HgCdTe array for T TTS). The SBC calculates the centroid position of the PSF on the TTS detector array and sends this information back to the observatory. The FSM in the telescope optical train⁷ corrects for the pointing accuracy and mitigates the long-term drift down to a few mas over a timescale of a few hours. The detector biasing and readout electronics need to be very stable so that the detectors themselves can achieve their relative photometric stability requirement, which will necessitate careful temperature control of the electronics boxes.

2.2.7.

Risk management approach

As an instrument on a NASA Class A mission, the MISC-T instrument has a fully redundant, dual-string, cross-strapped design. The MISC-T instrument has a warm electronics box each with two redundant sets of warm electronics, as shown in the block diagram (Fig. 1). Within an electronics box, the SBC and the rest of the warm electronics are cross-strapped as shown in Fig. 7. Although the focal plane arrays are not dual redundant, their multiplexer readouts are arranged so that sections of the arrays fail gracefully, as was done for the JWST instrument arrays.

Fig. 7

The cross-strapping design used in each of the MISC-T modules’ warm electronics provides dual-string fault tolerance.

2.2.8.

MISC-T alignment, integration, test, and calibration

The MISC-T module will require its own unique facilities for testing and calibration at the instrument level before it is delivered for integration with the observatory. Aside from the normal vibration, acoustic, and electromagnetic interference (EMI) environmental testing, the MISC-T module will need a thermal vacuum test chamber that can be cooled to liquid helium temperatures that is large enough to enclose the complete cold section of the MISC-T instrument and is IR light-tight enough that enough that stray light within the MISC-T detectors’ sensitivity wavelengths is no greater than the infrared background that would be expected on orbit. This test chamber should allow the introduction of IR signals at the expected levels of the on-orbit target sources and the ability to verify the spectral response of the spectrometer. Finally, the test facility will require an ultra-stable light source that can verify the ∼few ppm pixel-pixel detector stability requirements needed for the exoplanet transit measurements. This test facility will need to allow connections between the cold MISC-T instrument assembly in the test chamber and the warm MISC-T drive electronics over flight-like cables.

The warm MISC-T instrument drive electronics will not require any particular test facilities other than the usual test facilities used to conduct vibration, acoustic, EMI, and thermal vacuum tests. A spacecraft simulator will be required to test command and data transfers between the MISC-T and spacecraft electronics.

It will be necessary to develop specialized test hardware, but although the scope of this development will be significant, it is straightforward, with no new technology development required. MISC-T instrument I&T is expected to take 12 months, assuming all the external test support hardware and software is ready in time by the start of testing.

2.2.9.

MISC-T predicted performance

The MISC-T performance is calculated based on the following assumptions.

Throughput: The throughput curves as a function of the wavelength are calculated for T-S, T-M, and T-L, taking account of the reflectance of mirrors (four mirrors for telescope optics, three mirrors for fore-optics, three collimator mirrors, seven mirrors in the densified pupil spectrograph), the transmittance/reflectance of the dichroic beam splitters, grating efficiency, transmittance of the AR-coated lenses, detector quantum efficiency, and contamination/slit loss. The diffraction losses in the instrument were minimized with oversized optics (at least $12\lambda /D$). The final throughput including detector quantum efficiency is plotted in Fig. 8.

Fig. 8

The final throughput as a function of wavelength used for MISC-T performance calculations.

Zodiacal light: In addition to the light from the host star, there will be additional infrared light coming from the Zodiacal background. As the field stops have to be relatively large (3″ radius for T-S and T-M, 2″ radius for T-L) to capture the majority of the poorly focused, short-wavelength light from the $30\mu \mathrm{m}$ diffraction-limited Origins telescope, the shot noise from this extra light can be a significant source of signal stability errors at some wavelengths. The Glasse Zodiacal model¹⁰ is employed for the low background case outside the ecliptic plane while the Wright Zodiacal model for the ecliptic plane ($l=0$)¹¹ is employed for the high background case. The surface brightness of the zodiacal light, including scattered light and emissive components at representative wavelengths, is summarized in Table 5.

Table 5

The surface brightness of the zodiacal light models used for the MISC-T performance calculations.

Pointing jitter and long-term drift: The major source of high-frequency pointing jitter is due to the cryocooler and is expected to be $<50\mathrm{mas}$, with a current best estimate of 9 mas RMS. Low frequency pointing drift has a much stronger effect on spectral stability. The Lyot-Coronagraph-based TT sensor (TTS) in the MISC-T instrument is used for measuring the long-term drift so that it can be compensated with the telescope FSM, if necessary. Together the performance is predicted to be 1mas drift over a timescale of a few hours.

Detector performance: A $2\mathrm{k}\times 2\mathrm{k}$ HgCdTe detector array operated at 30 K is used in each of the MISC-T-S and T-M channels and a $2\mathrm{k}\times 2\mathrm{k}$ Si:As detector array operated at $\sim 8\mathrm{K}$ is used in MISC-T-L. The detector dark current is assumed to be $1{\mathrm{e}}^{-}{\mathrm{s}}^{-1}{\text{pixel}}^{-1}$ for the $2\mathrm{k}\times 2\mathrm{k}$ HgCdTe detector and $0.2{\mathrm{e}}^{-}{\mathrm{s}}^{-1}{\text{pixel}}^{-1}$ for the $2\mathrm{k}\times 2\mathrm{k}$ Si:As detector.¹² In addition to the science pixels, the reference and background pixels are applied for calibration of the systematic components present in the semiconductor device and ROIC. Since the time-variation component is common to all of the detector pixels, the common systematic components from the reference and background data are reconstructed and subtracted from the science data. The offset component estimated from the time- and pixel-averaged background data is also subtracted from the science data. This calibration procedures¹³ enable us to obtain the transit signal that is free from both the systematic and offset components. The dark current and background light are assumed to be constant during the transit observation. Thus, the stabilization of the detector’s temperature and calibration of the dark current value using the time series of the detector temperature are required.

Based on these assumptions, the team simulated the light curves of late-M type stars at different K-band magnitudes and calculated the spectro-photometric accuracy at representative wavelengths for each type of late-M type star. The results calculated for 3, 5, 8, 10, 14, and $20\mu \mathrm{m}$ assuming 84 transits and for a spectral resolving power of $R=50$ are shown in Figs 9Fig. 10Fig. 11Fig. 12Fig. 13Fig. 14Fig. 15–16. Noise sources due to photon shot noise, dark current, zodiacal light, jitter, and drift are included as shown. Aside from dark current, other detector and readout-electronics noise sources are not included. As an example, a bright star will have more integrated photons in 84 transits, so that the dark current noise contribution from the detectors will be relatively smaller. The flat contribution of the jitter and drift noise to the signal-to-noise over different host star brightnesses may seem odd at first, but in fact this term is a constant that is independent of the stellar magnitude.

Fig. 9

The current best estimate of the spectro-photometric stability achieved by MISC-T at $3\mu \mathrm{m}$ with $R=50$.

Fig. 10

Same as Fig. 9 but for at $5\mu \mathrm{m}$ with $R=50$.

Fig. 11

Same as Fig. 9 but for at $8\mu \mathrm{m}$ with $R=50$.

Fig. 12

Same as Fig. 9 but for at $10\mu \mathrm{m}$ with $R=50$.

Fig. 13

Same as Fig. 9 but for at $14\mu \mathrm{m}$ with $R=50$ (low background case¹⁰).

Fig. 14

Same as Fig. 9 but for at $14\mu \mathrm{m}$ with $R=50$ (high background case¹¹).

Fig. 15

Same as Fig. 9 but for at $20\mu \mathrm{m}$ with $R=50$ (low background case¹⁰).

Fig. 16

Same as Fig. 9 but for at $20\mu \mathrm{m}$ with $R=50$ (high background case¹¹).

3.2.1.

MISC WFI operation principle

The MISC WFI covers the wavelength from 5 to $28\mu \mathrm{m}$ with two channels; WFI-short (WFI-S) covers the wavelength from 5 to $9\mu \mathrm{m}$, and WFI-long (WFI-L) covers from 9 to $28\mu \mathrm{m}$. MISC WFI-S serves as the focal plane pointing and guiding for the observatory and has two redundant sets of $2\mathrm{k}\times 2\mathrm{k}$ Si:As detector arrays. During the normal operation, one of the two Si:As detector arrays in the WFI-S is used, and the mirror changer switches the optical path if needed to use the other detector in case the performance of the first detector is degraded for some reason. Since the WFI-S array is also critically needed for the MISC-T performance, it is important that it is redundant. It should be noted that between the two channels, WFI-L channel was prioritized for science promise over WFI-S by the Science and Technology Definition Team (STDT), in case a trade choice would need to be made between the two. The block diagram of MISC WFI is shown in Fig. 17.

Fig. 17

The block diagram of the MISC WFI.

3.2.2.

MISC WFI optical design

The MISC WFI offers a wide field imaging ($3^{\prime }\times 3^{\prime }$) and low-resolution spectroscopic capability with filters and grisms covering from 5 to $28\mu \mathrm{m}$ with a resolving power of $R\sim 300$. WFI-S and WFI-L share the same $3^{\prime }\times 3^{\prime }$ FOV by means of the beam splitter and the diffraction-limited image quality over the entire FOV is achieved at any wavelength from 5 to $28\mu \mathrm{m}$ with the help of a deformable mirror (DM) and a tip tilt mirror in the fore optics of the MISC WFI. The reflected beam (i.e., 5 to $9\mu \mathrm{m}$) by the beam splitter is diverted to WFI-S and the transmitted beam (i.e., 9 to $28\mu \mathrm{m}$) is sent to WFI-L. The MISC WFI optical design includes a A) wave-front error correction system in the fore optics and B) the imager optics (see Fig. 18).

Fig. 18

The MISC WFI optical design.

3.2.3.

MISC WFI detectors and read-out electronics

The MISC WFI-S optical path contains two $2\mathrm{k}\times 2\mathrm{k}$ Si:As detector arrays and the MISC WFI-L optical path contains a $2\mathrm{k}\times 2\mathrm{k}$ Si:As detector array, each of which is bonded to a Si readout multiplexer operating at $\sim 8\mathrm{K}$ to provide good detective quantum efficiency in the 5 to $28\mu \mathrm{m}$ wavelength range. This design has been used extensively in previous space missions such as Spitzer, AKARI, and JWST, and in the Stratospheric Observatory for Infrared Astronomy (SOFIA) instruments. The detector development work needed for the MISC WFI comprises a modest factor of two increase in the detector dimension format for the WFI module.⁸

Amplification of the signals from the detectors is a straightforward reuse of similar technology from previous space applications of Si:As BIB detectors such as Spitzer, JWST, and WISE, and will employ dedicated satellite chips that are located in close proximity to the detectors and also operate at cryogenic temperatures.

The readout electronics design of the MISC WFI detectors is very similar to that used in previous space missions and will incorporate flexible readout patterns and a variety of readout strategies including double-correlated and Fowler sampling technologies.

3.2.4.

MISC observing modes (upscope) and data rate

MISC WFI offers six observing modes (i.e., guiding, imaging, low-resolution slit spectroscopy, low-resolution slit-less spectroscopy, and two MIR scan mapping modes) in addition to the MISC-T mode. (See Table 7.)

Table 7

MISC observing modes (upscope) and data rates.

3.2.5.

MISC WFI thermal and mechanical design

MISC WFI instrument components are located in two thermal zones. The instrument optics and detectors are located in a cold zone at the same 4.5 K temperature as the telescope. The detectors have a weak thermal link to this cold temperature so they can be heated slightly and held at a stable operating temperature. This operating temperature is $\sim 8\mathrm{K}$ for the Si:As detector. The drive electronics that handle instrument data collection and communication with the spacecraft are located in the $\sim 300\mathrm{K}$ warm zone.

The MISC WFI solid model is shown in Fig. 19. Table 8 lists the mass breakdown of cold component estimated each for beryllium and cordierite models. To reduce MISC WFI instrument mass, we assumed beryllium as the baseline material for the mirrors, mirror support structures, and base plate.

Fig. 19

A 3D solid model of MISC WFI module.

Table 8

Mass breakdown of MISC WFI cold component.

The MISC WFI module has the following mechanisms: (1) an DM assembly in the fore optics of the WFI module, (2) a tip–tilt mirror (TTM) assembly in the fore optics of the WFI module, (3) a slit mask changer in WFI-S, (4) a slit mask changer in WFI-L, (5) a filter wheel for WFI-S, (6) a filter wheel for WFI-L, and (7) a mirror changer in WFI-S. Since the telescope optics is limited to diffraction-limited image quality performance only at $30\mu \mathrm{m}$ and longer wavelengths, the MISC WFI module needs to have its own internal wave-front error correction optics. The MISC WFI module includes a DM and a TTM in its fore-optics and achieves diffraction limited image performance at $>5\mu \mathrm{m}$ for sources within a wide $3^{\prime }\times 3^{\prime }$ FOV. The slit mask changer used in MISC WFI-S has two 4-in. slots (a slit mask and a hole) and that in MISC WFI-L has three 4-in. slots (two slit masks and a hole). The filter wheel assembly used in each of the MISC WFI-S and WFI-L has triple wheels. Each wheel has six 3-in. positions for band-pass filters and/or grisms for low-resolution spectroscopy. The mirror changer is placed just before the two Si:As detector arrays in the WFI-S to switches the optical path when needed to do so.

Table 9 lists the MISC WFI resource requirements (volume, mass, and power).

Table 9

MISC WFI module resource requirements.

3.2.6.

Instrument control

The MISCWFI has redundant dual-string warm electronics boxes (WEB A and WEB B for MISC WFI module). Each electronics box contains a SBC and four boards: (1) a CPI/ Power board, (2) a digital and analog board for array control (two $2\mathrm{k}\times 2\mathrm{k}$ Si:As arrays for WFI-S, one $2\mathrm{k}\times 2\mathrm{k}$ Si:As array for WFI-L), (3) a board for peripheral control for detectors (two $2\mathrm{k}\times 2\mathrm{k}$ Si:As arrays for WFI-S, one $2\mathrm{k}\times 2\mathrm{k}$ Si:As array for WFI-L), two slit mask wheels, two filter wheels, and one mirror changer, and (4) a control board for the cryogenic DM and TTM. A SBC is used to analyze the shape of the PSF and to give feedback to the DM and TTM to achieve the diffraction-limited image quality at $5\mu \mathrm{m}$. This SBC also provides reference star centroid information that is sent to the Origins telescope pointing control.

3.2.7.

Risk management approach

As an instrument on a NASA Class A mission, the MISC WFI module has a fully redundant, dual-string, cross-strapped design. This is particularly important for MISC WFI as the imaging array in the WFI-S may also serve as the focal plane pointing and guiding array for the observatory. Each of the MISC modules (i.e., MISC-T and MISC WFI) has its own warm electronics box each with two redundant sets of warm electronics. Within an electronics box the SBC and the rest of the warm electronics are cross-strapped as shown in Fig. 7. All of the moving mechanisms in the cold areas of MISC WFI have dual coils/actuators, with each coil/actuator capable of being driven by either of the A or B warm electronics sides. Although the focal plane arrays are not dual redundant, their multiplexer readouts are arranged so that sections of the arrays fail gracefully as was done for the JWST instrument arrays.

3.2.8.

MISC WFI alignment, integration, test, and calibration

Since the MISC WFI module is used by the observatory for pointing, this instrument must meet confocality requirements with all other Origins instruments. As these other instruments all operate in the far-infrared, meeting this requirement should not be difficult. The two-module configuration also greatly simplifies and speeds up instrument testing, as the modules can be tested independently and in parallel, if desired. Specialized test hardware development is required, but although the scope of this development will be significant, it is straightforward, with no new technology development required. The integration and test of the MISC instrument are expected to take 12 months, once all external test support hardware and software is ready.

3.2.9.

MISC WFI predicted performance

The following assumptions are made for the sensitivity estimate of the MISC WFI module (cf. Ref. 15). The $5\sigma$ 1-h line sensitivity of MISC WFI ($R=300$) is shown in Fig. 20.

Fig. 20

$5\sigma$ 1-h line sensitivity of OST/MISC WFI ($R=300$).

Exposure Time: The longest exposure time is determined as 300 s for the Si:As $2\mathrm{k}\times 2\mathrm{k}$ detectors ($30\mu \mathrm{m}$ pix-1; Raytheon) taking account of the number of pixels affected by cosmic ray hit events during an exposure and assuming the cosmic ray hit event rate at L2 as $5\times {10}^{-4}{\mathrm{m}}^{-2}{\mathrm{s}}^{-1}$.¹⁶

The Detector Performance: Detector Dark Current is assumed to be $0.06{\mathrm{e}}^{-}{\mathrm{s}}^{-1}{\text{pixel}}^{-1}$ for the Si:As $2\mathrm{k}\times 2\mathrm{k}$ detector.¹⁷ The readout noise (reduced up to 1/4 by means of Flowler-16 sampling) is assumed to be $10{\mathrm{e}}^{-}$ for the Si:As $2\mathrm{k}\times 2\mathrm{k}$ detectors.