MIRI Cross-Mode Recommended Strategies

A general set of guidelines can be applied to all MIRI observing modes.

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This page provides guidelines to aspects of the proposal planning process common to all MIRI observing modes: every program has to consider detector operations (i.e., how to choose a combination of groups and integrations per exposure), background strategies that will help to mitigate the spurious sky + thermal telescope emission, and general target acquisition (TA) aspects.

Recommended strategies for the main MIRI observing modes can be found in the following JDox pages: MIRI Imaging Recommended StrategiesMIRI MRS Recommended StrategiesMIRI LRS Recommended Strategies, MIRI Coronagraphic Recommended Strategies, and MIRI TSO Recommended Strategies.

Detector readout recommended strategies

See also: Understanding Exposure Times, MIRI Detector Readout

Like other instruments on JWST, MIRI detectors use MULTIACCUM readouts of multiple groups along the integration ramp. Once the final group in an integration is read, the detector circuit is immediately reset and, if defined, a new integration starts. The on-sky time of each individual exposure (i.e., the time spent in a single dither position) is defined by the number of groups and integrations.

How many groups and integrations should I use?

For MIRI the optimal combination of groups and integrations depends on the target brightness, the background and the desired signal-to-noise ratio.

Figure 1. Flow diagram illustrating general recommendations for the MIRI detectors usage

For time-series observations, please see MIRI TSO Recommended Strategies.

What is the recommended ideal, minimum, and maximum length of an integration?

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The MIRI Si:As detectors exhibit a set of non-ideal, but known, detector effects (see Ressler et al. 2015). The first few groups in an integration are most strongly affected. An increasing number of groups per integration can mitigate these effects.  

  • Ideal number of groups will depend on the source brightness and background conditions:
    • Bright sources/high background conditions will generally mean that saturation is reached in a short number of groups (5 to 10). Having one group saturating at the end of the integration will maximize the dynamic range. For very bright cases, imaging users should favor subarrays.

    • Faint source observations will benefit from having long integrations. In FAST mode, 100 groups (about 280 s in FULL array, shorter when using subarray mode) is suggested as a starting point. Increase or decrease the integration to achieve your desired signal to noise. In SLOW mode, 25 groups (about 600 s) is a good starting point. Note that SLOW mode is recommended primarily to limit data volume when MIRI is being used in parallel.

  • Minimum number of groups: Five groups is typically the minimum number required to obtain a reasonable calibration. For very bright sources, observers are permitted to use less than 5 groups, but there is no  guideline for the noise and photometric accuracy we can expect in these very short integrations. When using fewer than 5 groups, and if photometric accuracy is important to your program, you should plan to observe a calibration star with exactly the same exposure parameters. If, instead of accuracy, repeatable precision is your primary need (e.g., transits), then this step may not be necessary.

  • Maximum recommended integration length: 360 groups in FAST/FULL mode, 42 in SLOW mode (approximately 1,000 s). This is not a hard limit, it is a recommendation based on a study that considers cosmic ray rates. The study has shown that after a 1,000 s integration, most of the detector pixels are expected to be affected by cosmic rays. This not only includes pixels that are directly impacted by a cosmic ray, but also its adjacent pixels. Note that the ETC will model these results using a statistical approach by which each cosmic ray affects 9 pixels (see Figure 2).

Users are strongly encouraged to follow this general guidelines of integration length optimization.

Figure 2. Pixel affected by a direct cosmic ray hit

Left: The illustration depicts a jump in a pixel integration after receiving a direct cosmic hit, and the secondary and tertiary jumps on the 8 adjacent pixels. This illustrates the statistical approach used by the JWST Exposure Time Calculator (ETC) that assumes each cosmic ray affects a total of 9 pixels (pipeline). Note the effect in the corner pixels (tertiary jump) is smaller than in the adjacent (secondary) ones.
Right: Illustration of the pixel single integration exposures affected by the cosmic ray. Top Right: Primary jump on the central pixel. Bottom Right: Secondary jump.

How should I deal with saturation?

In many cases, the maximum integration length will be set by saturation of bright sources within the field, emission lines within a spectrum, or background emission. Because MIRI samples many points "up the ramp", saturating midway in the integration does not mean that data are lost; the pipeline will only fit the initial unsaturated data points. This can be an advantage since you can integrate longer to gain better signal to noise on fainter areas of the field or spectrum. However, saturation, as well as bright sources, will leave latent images in subsequent exposures (although they do not appear to be a strong function of the integrated signal level). High redundancy in the data (e.g., dithering) is recommended to mitigate the effects of persistence.

Should I use multiple integrations?

The first integration within an exposure differs from the second and subsequent ones. This is because there are multiple detector resets between exposures (the system goes into a continuous reset while the telescope is dithering, filter wheels are moving, etc.), allowing some of the above-mentioned transient features to be cleared out of the first integration of the exposure. Within a single exposure, however, there is only a single reset between integrations; this leads to a flux-dependent decay effect for the first few groups of each successive integration. In general, the longer the integration, the more accurate the result (i.e., the errors in the correction become less significant with more samples).

To decide whether to use multiple integrations, observers should consider various aspects: brightness of the source, detector performance, dithers, and overheads. These can all be encapsulated in 2 broad cases:

  1. Bright sources/high backgrounds that saturate rapidly: in this case, multiple integrations that allow at least one saturated group will be beneficial to obtain high SNR, and provide better efficiency. Telescope maneuvers to the next dither position are more "costly" than starting a new integration.

  2. Fainter sources that can reach the "ideal number" of 100 groups in FAST mode and 25 in SLOW (see above) will benefit from single integrations in all dither positions. It is best to select integration lengths which are as long as your observations allow (i.e., a few long integrations are better than many shorter ones).

Choosing the readout mode: FAST vs. SLOW

See also: MIRI Detector Readout Overview

The MIRI detectors can be operated using two different readout modes: FAST and SLOW. The main differences between these modes are:

  • FAST mode reads out the detector every 2.775 s in full array mode, vs. 23.889 s in SLOW mode. FAST mode offers finer time sampling that allows better characterization of detector effects and more samples for cosmic ray correction. Compared to SLOW mode, in FAST mode there is approximately a factor of 9 less time loss in case of a cosmic ray hit.

  • SLOW mode provides about 9 times less data volume.

FAST is the recommended readout mode for all MIRI observations executed as prime. SLOW mode is useful for parallel observations, where it can be essential to limit the data volume.

 Background observations recommendations

See also: JWST Background Model, JWST Background-Limited ObservationsMRS Dedicated Sky Observations

Does my program need background observations? 

Observers should carefully consider the impact the extra emission of the JWST background's (modeled by the JWST ETC Backgrounds) will have on their data. The recommended strategy to account for and remove the background depends on the source:

  • Point sources: For point source studies, observers do not need to include additional background observations; the region off-source pixels can be used to derive the background. In high background observations, the flat field uncertainty will affect the SNR.

  • Extended sources: When the emission of the science target covers the entire FOV, it is advised to obtain background data in every spectral configuration used by the science data. This can be done in 2 different ways.
    • Ideally the background region should be within your mosaic. This is because the background matching step in the pipeline will match the fluxes in the overlapping regions; the net result is that the photometric zeropoint will be lost (i.e., users won't accurately know where the zero flux/background level is). With this strategy the mosaic contains a region with only background contributions.
      • The calibration pipeline will not apply background corrections using a region contained in a mosaic.
    • Sometimes, when the source is very large, this will not be feasible. In that case users can define a separate background target (see details in APT Targets). It is advisable, then, to verify the individual exposures (before they are combined by the pipeline) to assess the individual photometric zeropoint on each of them.
      • The calibration pipeline will stack all dithered background images (to remove sources and artifacts) and subtract the result from the science exposures.

To understand whether the Background Limited special requirement in APT should be used, please read this JWST Background-Limited Observations article.

How often do I need to get a background?

The JWST mission defines visits as individual schedulable units, used to build up the observation timeline. Visits that are not linked by special requirements can be planned at different times of the observing cycle and background variations are expected. Observers should plan on taking background data for every observing period (i.e., non-linked visit) in their programs. The zodiacal background will vary with a timescale of weeks. 

Pre-launch, it is difficult to predict the degree to which the telescope self-emission will vary temporally. Observing programs that need low backgrounds can request visits to be scheduled when the background is predicted to be relatively low. Observing programs that require high accuracy relative calibrations in the results (better than 1%) and use extended targets may consider taking background data before and after the science exposures. Users are encouraged to use the JWST JWST Backgrounds Tool to better understand the impact of the background in their observations, its intensity, and components as a function of time.

Target acquisition

An overview of the MIRI Target Acquisition process is given elsewhere in the documentation and discussed in the mode-dedicated target acquisition pages (MIRI Imaging Target Acquisition, MIRI MRS Target Acquisition, MIRI LRS Slit Target Acquisition, and MIRI LRS Slit Target Acquisition). These articles also include guidelines on when TA is needed for each MIRI mode.

TA targets

See also: MIRI Target Acquisition

When TA is needed (see JWST Pointing Performance) the science target is typically used for TA. However, the procedure can also be carried out with a nearby bright star, which should be within the APT visit splitting from the science target. The splitting distance ranges between 30"–80", depending on the Galactic latitude of the target, and is defined as the total distance between the TA and science exposures which is the combination of:

  1. The distance between the TA target and science target.
  2. The distance between the TA aperture and the science aperture.

If the TA target is not within this visit splitting distance, the observation will not be schedulable by the APT Visit Planner. Users are encouraged to check whether their program can be scheduled when using off-source TA targets.

If feasible, using an offset TA target should be considered in the following scenarios:

  • The science target is spatially resolved, resulting in a higher uncertainty on the centroid location (see section below).
  • The science target's spectral energy distribution is not well know in the MIRI TA filters, leading to an uncertain estimation of the exposure time and SNR.
  • The science target requires a long (~100 s or longer) integration to reach SNR of 20. 

Understanding the TA onboard procedure

The aim of this section is to give details on the onboard TA data process, so users understand the several aspects that might impact its accuracy/outcome. The onboard centroid algorithm for MIRI works as follows:

  • The onboard algorithm uses the first, middle and one-before-last frame to generate two 2-frame difference images. The final image used for TA is then constructed by taking the minimum value of these 2 difference images on a pixel by pixel basis.
  • This raw image is pre-treated (background subtracted and flat fielded).
  • It then finds the brightest 3 × 3 pixel region of the detector region of interest (ROI, see Table 1). The checkbox size (3 × 3) has been defined to encompass the imager PSF.
  • After that, it performs a fine location by calculating the center of mass in the previously located brightest area.

The accuracy of this procedure depends on several aspects:

  • The TA target should be a point source. The algorithm will work in extended/resolved sources, but with less accuracy.
  • The integrated signal-to-noise ratio should be ~20.
  • The TA source should be the brightest source in the ROI. If there is a source brighter (by any factor) than the one selected by the observer, TA will be carried out on that. If there is a source with the same brightness the algorithm will perform TA in the first one encountered following the direction in which the detector is read out.
  • Users should be aware that regions with bright diffuse emission may also result in false identification of the TA target. To reduce the possibility of this happening, the checkbox area is a very small one. 

Table 1. Sizes of detector regions of interest used to perform target acquisition. 

TA region of interest (ROI)Size
MRS (TA performed in imager detector)48 × 48 pixel2
LRS (slit and slitless)48 × 48 pixel2
Coronagraphs48 × 48 pixel2 (16 × 16 pixel2 secondary TA)

Note that the MIRI detector plate scale is 0.11 arcsec/pix.

See more details on the MIRI Imaging Target Acquisition, MRS MIRI MRS Target Acquisition, MIRI LRS Slit Target Acquisition, and MIRI LRS Slit Target Acquisition articles.

Figure 3. Depiction of a dense MRS TA region (48 × 48 pixels in the MIRI imager)

In this example of a rich field, if the observer selects the central star as a TA target, the centroid algorithm will not be performed there but in the brightest pixel of the field (top right). After the TA is performed on the wrong target, data will not be taken at the coordinates of the science target specified in the proposal. Observers are encouraged to carefully select the TA target so that there are no brighter neighbors.

Target acquisition readout mode: FAST and fast group averaging

MIRI target acquisition can be performed using 6 different readout modes: FAST and 5 different flavors of fast group averaging (see MIRI Target Acquisition). As data volume can become a concern when TA requires a large number of groups, in fast group averaging each group consists of 4, 8, 16, 34, or 64 co-added FAST mode groups. Fast group averaging is offered for MIRI TA only. The ETC will give warnings when the SNR achieved is not sufficient to successfully finish the TA procedure (see more details in JWST ETC MIRI Target Acquisition).



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Stray light field dependence for the James Webb Space Telescope

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