NASA

Roman Space Telescope

Wide-Field Instrument : Technical

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The Wide-Field Instrument (WFI) is the primary instrument of the Roman Space Telescope. WFI provides wide-field imaging and multi-object, slitless spectroscopy over an. 0.281 deg2 field-of-view. The design of WFI incorporates an 11-position element wheel assembly, 18 optical/NIR sensitive detectors mounted to a mosaic plate assembly and linked to associated control electronics, and a relative calibration system. The WFI detectors also acquire guide star images interleaved with science data collection to enable fine pointing control. The key parameters of the WFI are summarized below.



INSTRUMENT PARAMETERS

Field of View

0.8 x 0.4 deg, 0.281 deg2 active area (excluding detector gaps)

Spatial sampling

0.11 arcsec/pix; 288 Mpix focal plane

Image Stability

1.0 nm RMS wave front error variation in 180 sec

Detector Mosaic

6 x 3 array of Teledyne H4RG-10

Band pass

0.48 to 2.3 microns

Element Wheel positions

8 science filters, grism, prism, and blank for dark current and flat fields (using the internal calibration system)

Instrument + Telescope Focal Ratio

f/7.9

Engineering Test Unit detector focal plane

Image of Engineering Test Unit detector focal plane

The Roman WFI detector focal plane consists of 18 H4RG-10 detectors in a 6 x 3 array that combined provide an 0.281 deg2 active FOV. This image shows the focal plane Engineering Test Unit with non-flight candidate detectors.



Filters and Imaging

WFI carries 8 science filters with overlapping band passes spanning 0.48 – 2.3 microns. The filters are located at the telescope exit pupil. They are housed in the element wheel and are rotated into the instrument light path for multiwavelength imaging.

Filter Point Spread Functions

Point spread functions (PSFs) for the Nancy Grace Roman Space Telescope have been created using WebbPSF, a python based package. This tool takes into account properties of the telescope and the instruments, including detector pixel scale, rotations, filter profiles, and point source spectra. These are not full optical models, simply a tool that transforms the optical path difference maps, into the resulting Roman PSFs.

The website linked above provides instructions on how to install WebbPSF, how to run it via the Python API, in addition to providing Roman specific examples.

A small subset of simulated Roman PSFs for a K0V star can be found here (wim_psf_subset.zip - 28MB). These PSFs were calculated at the centers of each SCA (1 through 18), and at the boundaries of the FOV. An 8 mas pointing jitter in each axis was included in the PRF formulation, and each PRF was oversampled by a factor of eight.

More extensive libraries, with PSFs computed at the center of each SCA and at various locations around the boundary of the FOV are also available; these have been computed with stellar spectral energy distributions for spectral types:

Filter parameters


Element nameMin (μm)Max (μm)Center (μm)Width (μm)RPSF FWHM (arcsec) *
F0620.480.760.6200.2802.20.058
F0870.760.9770.8690.21740.073
F1060.9271.1921.0600.26540.087
F1291.1311.4541.2930.32340.105
F1581.3801.7741.5770.39440.127
F1841.6832.0001.8420.3175.810.151
F2131.952.302.1250.356.070.175
F1460.9272.0001.4641.0301.420.105

* Note: PSF FWHM in arcseconds simulated at the center of the center pixel of a detector near the center of the WFI FOV using an input spectrum for a K0V type star. Please click on the FWHM value for each filter to view the simulated PSF.

The above table provides representative PSF FWHM values for a K0V type star near the center of the WFI FOV. Simulated PSFs for different stellar types and all detectors across the focal plane are available here.

Filter Effective Area Curves

The effective area as a function of wavelength for each Roman filter is available in tabular form here (.XLS).

Imaging Sensitivity Overview

The table below gives AB magnitude limits achieved at 5σ in 1 hour or 55 seconds of integration, for twice the minimum zodiacal light background, both for point sources and for compact galaxies with half-light radius of 0.3”. In the one-hour cases, the SNR calculations assume co-addition of four 900s exposures, while the 55s case assumes a single exposure. Profile fitting photometry is assumed.

FilterF062F087F106F129F158F184F213F146
Wavelength (microns)0.48-0.760.76-0.980.93-1.191.13-1.451.38-1.771.68-2.001.95-2.300.93-2.00
1 hr, Point28.528.228.128.028.027.426.228.4
1 hr, r50=0.3”27.126.826.826.826.826.425.227.3
55s, Point25.525.125.125.024.924.423.725.9
55s, r50=0.3”24.123.823.823.823.823.322.724.7

Scalings for other deep cases: The flux density limit for other deep integrations, between about ten minutes and tens of hours, can be estimated from the 1 hour depths using a flux t-1/2 scaling. For integrations below about 10 minutes, this scaling becomes over-optimistic by > 30% due to read noise and overheads. Limits for a half-light radius of 0.2” are slightly closer to the 0.3” case than the point source case. For more extended sources, the limiting flux density fν,limfor fixed SNR and integration time should be scaled from the 0.3” case approximately as fν,limr50 or ΔABlim=2.5 log⁡(0.3"/r50). Reducing the assumed zodiacal background from 2x to 1.44x minimum improves deep imaging flux limits by about 0.15 mag for the F062 through F158 filters, 0.08 mag for F184 and F146, and negligibly for F213.

Fast/Wide Limit: The final two lines in the table give magnitude limits achieved at 5σ in 55 seconds (with a single exposure). At this integration time, Roman can cover approximately 8 contiguous square degrees per hour in one spectral element, and slew-and-settle overheads slightly exceed integration time. There is little point considering faster survey speeds, because sensitivity drops rapidly for modest increases in survey speed at yet shorter integrations.

More information on the filter sensitivities is available in tabular format on the Anticipated Performance Tables page.


Roman Effective Area

[View PDF of Filter effective area vs. wavelength]
The effective area as a function of wavelength for each Roman filter is available in tabular form here (.XLS).

Zodiacal Light and Thermal Background

The tables below provides the count rate per pixel at minimum Zodiacal light in each filter and the estimated thermal background. For observations at high galactic latitudes, the Zodi intensity is typically ~1.5x the minimum. For observation into the galactic bulge the Zodi intensity is typically 2.5-7x the minimum.

Count rate per pixel at minimum Zodiacal Light
F062F087F106F129F158F184F213F146
0.270.270.290.280.260.150.130.85
Internal thermal backgrounds (count rate per pixel)
F062F087F106F129F158F184F213F146
00000.040.174.520.98


Imaging Sensitivity Calculator

On this page you will find tools designed to help you, the user, calculate the exposure time required for a given source at a given signal to noise, or vice versa.

We have provided a Jupyter notebook which walks through each step of the calculation, and a python script that will open a gui in which you can input your objects information.

To run both scripts you will need to download the accompanying data files.

Both the notebook and the gui require the user to specify the filter, zodiacal light contribution, type of source, fitting method, and signal to noise. In this notebook we will be doing exposure time calculations for point sources, and extended sources with half-light radii of 0.2 arcsec or 0.3 arcsec.

The exposure times used within these calculations are quantized in multiples of 3 readout frames, with the number of visits/dithers being 1.

An example of using the Jupyter notebook is provided here.




Grism and Prism

WFI carries 2 dispersive elements for slitless, multi-object spectroscopy. The grism band pass spans 1.0 – 1.93 microns and has a resolution of ~600. The prism band pass spans 0.75 – 1.80 microns and has a resolution of ~100. The dispersing elements are housed in the element wheel and are rotated into the instrument light path for slitless spectroscopy across the WFI FOV.

Grism and Prism parameters


Element nameMin (μm)Max (μm) Center (μm)Width (μm) R
G1501.01.931.465 0.930461λ(2pix)
P1270.751.801.2751.0580-180 (2pix)

Grism and Prism Effective Area Curves

The effective area as a function of wavelength for the Grism and Prism are available in tabular form here (.XLS).

Grism and Prism effective area vs. wavelength

The effective area as a function of wavelength for the Grism and Prism are available in tabular form here (.XLS).

Grism and Prism zodiacal light

The table below provides the count rate per pixel at minimum zodiacal light for the grism and prism. For observations at high galactic latitudes, the Zodi intensity is typically ~1.5x the minimum. For observation into the galactic bulge the Zodi intensity is typically 2.5-7x the minimum.


Count rate per pixel at minimum zodiacal light
GrismPrism
0.650.95

Grism Spectroscopy Sensitivity

The Roman WFI slitless grism has a spectral range of 1.00-1.93 microns and a dispersion of about 1.1 nm/pixel, essentially independent of wavelength, yielding a 2-pixel resolving power of R = λ / δλ = 460 λ / μm for a point source. The table below gives representative 1-hour sensitivities for both line and continuum. The continuum limit assumes coaddition of 2 pixels (2.2 nm) in the dispersion direction.

5σ limits for Roman WFI grism in 1 hour on source, 2x minimum zodiacal background.
Emission line limits are in units of 10-17 erg cm-2 s-1 , and continuum limits are in AB mags for 2 pixels
Wavelength (microns)1.051.11.21.31.41.51.61.71.81.9
fline,17, r50=0; 1 hour5.54.43.32.92.82.833.23.84.2
mAB(2 pix), r50=0; 1 hour21.421.621.721.721.621.421.221.020.720.5
fline,17, r50=0.3"; 1 hour14.411.48.57.36.97.07.47.48.69.5
mAB(2 pix), r50=0.3",1 hour20.420.620.720.720.620.4 20.220.119.819.6

Sensitivities for other integration times (between a few minutes and tens of hours) and zodiacal backgrounds can be scaled from the above using f limt-1/2 b1/2, where t is integration time and b the zodiacal background level. Because this is slitless spectroscopy, the scaling of sensitivity with size for rr50 > 0.3" differs between line and continuum sensitivity, with limiting behaviors of flne r50, and fcont r501/2.

Prism Spectroscopy Sensitivity

The Roman WFI slitless prism has a spectral range of 0.75-1.80 microns and a resolution that is strongly wavelength dependent, with 80 < λ / δλ < 180. The highest resolution is at the blue end of the prism wavelength coverage. In addition to its lower dispersion, the prism has higher throughput than the grism, making it more sensitive to continuum.

5σlimits for Roman WFI prism, 2x minimum zodiacal background
Wavelength0.801.001.201.401.601.75
Δλ (for 2 pixels, in nm)4.48.712.616.518.418.6
mAB(2 pix), r50=0; 1 hour23.023.523.823.823.823.7
mAB(2 pix), r50=0.3"; 1 hour22.422.923.223.223.223.1
mAB(2 pix), r50=0; 55 sec20.420.921.121.121.121.0
mAB(2 pix), r50=0.3"; 55 sec19.820.320.520.520.520.4

The same scalings that apply to grism spectroscopy can be used to scale other deep prism sensitivities from the 1 hour case.

More information on the grism and prism sensitivity is available in tabular format on the Anticipated Performance Tables page

Grism and Prism dispersion

The grism has constant dispersion and linearly increasing resolving power. The prism provides higher throughput and lower dispersion than the grism. The prism dispersion varies with wavelength and varies slightly with field angle.

Additional information on the dispersion of the grism and prism in tabular form is available here. (.XLS)

Prism dispersion curve


Detectors

The WFI detectors are Teledyne H4RG-10 models with 10 micron pixel pitch. Each of the 18 detectors has 4096 x 4096 pixels and 32 readout channels. The detectors have very low dark current at their 95 K operating temperature and low noise. Detailed information on the design and properties of the Roman detectors is available in Mosby et al. 2020. The key parameters are summarized below.

ParameterValueUnitsComment
Detector TypeTeledyne H4RG-10N/AN/A
Total pixels4096 x 4096pixelsRows x Columns
#Active (Columns)4088pixels4 reference +
4088 light sensitive +
4 reference
#Active (Rows)4088pixels4 reference +
4088 light sensitive +
4 reference
Pixel Pitch10 micronsN/A
Plate Scale0.11arcsecN/A
Readout Channels32N/AN/A
Readout rate200 kHzN/AN/A
Readout time3.04sN/A
Cutoff Wavelength2.5micronsN/A

Detector measured quantities

 
QuantitySpecification ValueTypical actualUnits
Dark CurrentDark current is the leakage current associated with the detector material when they are not being exposed to external radiation.
<0.1<0.005e-/sec/pixel
CDS Read NoiseCDS noise is the RMS noise from the difference of successive detector readouts.
<2015-16e- RMS
Total Noise in 180sTotal noise is the uncertainty on the slope fit for a sample-up-the-ramp sequence for an exposure time of 180 sec.
5-6e- RMS
Quantum EfficiencyQuantum efficiency quantifies the fraction of incident photons that get detected as signal. Quantum efficiency varies as function of wavelength due to the specific detector material, coatings, and pixel design.
>80%~95%(avg 0.8-2.1μm)
Quantum EfficiencyQuantum efficiency quantifies the fraction of incident photons that get detected as signal. Quantum efficiency varies as function of wavelength due to the specific detector material, coatings, and pixel design.
>60%75-80%(avg 0.6-0.8μm)
Pixel operabilityPixel operability refers to the number of detector pixels meeting all specifications.
>95%>99%N/A
PersistencePersistence is the observed release of photogenerated charge from a previous readout in a subsequent readout. It appears as a memory effect in the detector.
 ~0.1 e-/s/pixel








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