The Dark Energy Survey

Science Requirements and Technical Specifications

**Version 6.3 Draft**

*James Annis*

*Stephen Kent*

*Huan Lin*

*Tim Abbott*

*Brenna Flaugher*

*Wyatt Merritt*

12 July 2006

Contents

[1 Science Goals 4](#science-goals)

[1.1 The science objective 5](#the-science-objective)

[1.1.1 DECAM as a Community Instrument
5](#decam-as-a-community-instrument)

[1.2 Strategy of the survey 6](#strategy-of-the-survey)

[1.3 Science data products and science verification
6](#science-data-products-and-science-verification)

[2 Science Requirements 6](#science-requirements)

[2.1 Survey area and survey time span
7](#survey-area-and-survey-time-span)

[2.2 Wavelength coverage 8](#wavelength-coverage)

[2.3 Limiting magnitudes 9](#limiting-magnitudes)

[2.4 Point spread function 11](#point-spread-function)

[2.5 Throughput 12](#throughput)

[2.6 Photometric redshifts 13](#photometric-redshifts)

[2.7 Photometric calibrations 14](#photometric-calibrations)

[2.7.1 Linearity: 15](#linearity)

[2.7.2 Ghost subtraction: 15](#ghost-subtraction)

[2.7.3 Flatfielding: 16](#flatfielding)

[2.7.4 Residual stray light subtraction:
16](#residual-stray-light-subtraction)

[2.7.5 Variations in aperture correction:
16](#variations-in-aperture-correction)

[2.7.6 Primary standard star calibration:
16](#primary-standard-star-calibration)

[2.8 Astrometric calibration 17](#astrometric-calibration)

[2.9 Mapping 18](#mapping)

[2.10 Summary of Science Requirements
20](#summary-of-science-requirements)

[2.11 Success criteria 21](#success-criteria)

[3 Technical Specifications 23](#technical-specifications)

[3.1 Survey Footprint 23](#survey-footprint)

[3.2 Field of view 23](#field-of-view)

[3.3 Pixel scale 25](#pixel-scale)

[3.4 Point spread function 26](#point-spread-function-1)

[3.5 Radioactivity and Charged Particle Events
28](#radioactivity-and-charged-particle-events)

[3.6 Science CCD Specifications 28](#science-ccd-specifications)

[3.6.1 Quantum efficiency 28](#quantum-efficiency)

[3.6.2 QE and Temperature Stability 29](#qe-and-temperature-stability)

[3.6.3 Read noise 29](#read-noise)

[3.6.4 CCD Diffusion and Flatness 30](#ccd-diffusion-and-flatness)

[3.6.5 Cosmetic defects 31](#cosmetic-defects)

[3.7 Telescope tracking, focus and active optics
31](#telescope-tracking-focus-and-active-optics)

[3.7.1 Guide CCDs 32](#guide-ccds)

[3.7.2 Focus & alignment CCDs 32](#focus-alignment-ccds)

[3.8 System Requirements for the DES optical corrector
33](#system-requirements-for-the-des-optical-corrector)

[3.8.1 Passbands: 33](#passbands)

[3.8.2 Ghosting 34](#ghosting)

[3.8.3 Focal plane flatness 35](#focal-plane-flatness)

[3.8.4 Operating temperature range 35](#operating-temperature-range)

[3.8.5 Physical Size 35](#physical-size)

[3.8.6 Glass Type 35](#glass-type)

[3.8.7 Throughput 36](#throughput-1)

[3.8.8 Filters 36](#filters)

[3.8.9 Filter transmission 36](#filter-transmission)

[3.8.10 Dewar window 37](#dewar-window)

[3.8.11 Shutter 37](#shutter)

[3.8.12 Cost 38](#cost)

[3.9 Instrument Cage Motion 38](#instrument-cage-motion)

[3.9.1 Focus 39](#focus)

[3.9.2 Tilt and X&Y translation 39](#tilt-and-xy-translation)

[3.9.3 Orientation changes 40](#orientation-changes)

[3.10 Summary of Technical Specifications
41](#summary-of-technical-specifications)

# Science Goals

Our principal science goal is to measure ω, the dark energy equation of
state parameter, to a precision of δw ≤ 5% (1 σ) in each of four
separate and independent techniques. At the end of the experiment we
will combine into a complementary constraint with higher precision, but
we will also be able to cross compare the results for signs of differing
systematic errors.

The Dark Energy Survey is divided into four parts:

1)  **Cluster Surveys**: The redshift distribution, the spatial
    distribution, and the evolution of the mass function of clusters of
    galaxies are sensitive to w through standard volume and standard
    fluctuation probes. To accomplish measurements of these, the DES
    instrument must be able to locate and measure the redshifts of
    clusters of galaxies out to z ≈ 1, beyond the z where dark energy
    becomes important in the expansion history. The survey must cover
    large solid angle to enable the spatial power spectrum measurement
    and to deliver large cluster samples. The survey must have multiple
    bandpasses for precision photometric redshifts, and faint limiting
    magnitudes to increase the number of galaxy photometric redshifts
    available for each cluster. Finally, and very importantly, the
    survey must cover the same area as the South Pole Telescope
    collaboration Sunyaev-Zeldovich (SZ) survey so that the
    collaborations can share position, redshift, and SZ decrements,
    increasing the power of both surveys.

2)  **Weak Lensing**: The weak lensing of background galaxies by
    foreground mass sheets, correlated with itself (shear-shear
    correlations) or with foreground galaxies (shear-galaxy
    correlations) is sensitive to w as standard ruler and growth
    function probes. The DES Instrument must be able to measure the weak
    lensing signal from large scale structure in several redshift bins
    with small and stable instrumental point spread functions. Large
    solid angle is required to probe the shear field over the largest
    possible range in physical scales, and a small and stable PSF is
    required to make maximum use of the galaxies detected and to
    minimize the loss of efficiency due to systematic errors.

3)  **Galaxy Angular Power Spectra**: The spatial angular power spectrum
    of galaxies as a function of redshift is sensitive to w through a
    standard ruler probe. To accomplish this measurement, we must have a
    stable photometric calibration, cover large solid angle to probe a
    wide range of physical scales, and have faint limiting magnitudes to
    increase our sample size and probe to redshifts beyond z ≈ 1.

4)  **Supernovae**: The apparent magnitude of Type Ia supernovae
    provides a standardizable candle probe. The DES Instrument will be
    used to measure light curves of 2000 supernovae out to z = 0.8. The
    observations must cover a moderately large solid angle repeatedly to
    obtain light curves of the SNe. These SN observations must be taken
    in multiple bandpasses to provide colors of the objects, both for
    reddening measurements and SN type discrimination based on colors.
    Excellent red response is necessary to obtain adequate signal for
    higher redshift SN light curves. Detailed knowledge of the system
    passbands as a function of wavelength is required for accurate k
    correction calculations.

## The science objective

These science goals will be achieved by conducting a multicolor imaging
survey sufficient to create a deep galaxy map. The design of the survey
enables production of cluster catalogs 90% complete and pure out to z =
1, with a point spread function clean and stable enough to provide a
weak lensing measurements of large scale structure, photometric
calibration stable enough to allow the angular correlation function
measurements, and red sensitivity and time domain coverage sufficient to
allow z ≈ 0.75 supernova to be studied. These surveys collectively are
called the primary survey, and serve as the unifying science objective
of the collaboration.

The primary survey is the combination of the imaging survey, the time
domain survey, and the calibration data necessary to calibrate the
surveys. The primary survey is to be finished in a span of 5 years using
30% of the available observing time.

*The Imaging Survey*

The imaging survey consists of:

i.  a CCD survey of 5000 sq-degrees predominantly at Dec ≤30°;

ii. deep enough to detect 1/2 L\* galaxies out to z=1 in 4 bands (*g*=
    24, *r*= 24, *i*=24, *z*=23.6 10σ, for small galaxies);

iii. photometric redshifts for complete samples of galaxies to a density
     of roughly 20/sq-arcminute;

iv. data in a form that enables state-of-the-art studies of the large
    scale distribution of clusters, galaxies, and the weak lensing
    induced shear.

*The Time Domain Survey*

The time domain survey consists of:

v.  a CCD survey of roughly 40 sq-degrees in 13 fields sampled every
    three nights in *r* and every 6 nights in *i* and *z*.

vi. 1400 supernova light curves over 0.25 ≤ z ≤ 0.75, with photometric
    redshifts for the host galaxies.

### DECAM as a Community Instrument

The Dark Energy Survey instrument will be a community resource.
Community data taken with it using the standard observing mode form an
auxiliary resource for our science programs. We have been encouraged to
provide \"reasonable\" UV response or at least to not preclude the
ability to make UV measurements.

## Strategy of the survey

Our approach to conducting the imaging survey starts with placing a
single large camera onto a single large telescope. We then use the
instrument to image the survey area with multiple overlapping tilings at
a cadence of two tilings per year per bandpass, on average, using the
overlapping data to perform survey wide relative photometry, tie down
the absolute calibration by subsampling tiles for standard star transfer
observations, and finally to coadd the tilings to obtain the depth
necessary to measure z ≈ 1 clusters.

The coadded tiling is searched for stars and galaxies, these are
measured, and the results grouped into and star and galaxy catalogs. The
galaxy catalogs are supplemented by photometric redshift estimates.
These form the primary data products collaboration scientists use.

## Science data products and science verification

The deliverables from the survey will consist of the science data
products, such as object catalogs and coadded images, and the core set
of science analyses for the four key projects.

The galaxy catalog, with associated photometric redshifts and star
catalog, will provide the basis for a set of derived science data
products. These will include a cluster catalog, a weak-lensing optimized
shape measurement catalog, and a time domain catalog. A product from
outside the collaboration that will be available to the collaboration is
the SPT SZ catalog.

Using these science data products, the science teams will conduct to the
actual science analyses of cluster counting, galaxy angular power
spectra, weak lensing cross correlation tomography, and supernova light
curves, and derive estimate of *w* and other cosmological parameters.

In addition to the science data end products, there will be graphs,
analyses, and correlations that the team will create with the data that
will be used to determine that the science requirements and goals have
been achieved.

# Science Requirements

The form of the science requirements is important. The aim is to lay
quantitative requirements down on the fundamental quantities that must
be measured and controlled to have a successful experiment. We adopt a
model from the Sloan Digital Sky Survey. For each requirement there is a
list of questions to be answered:

i.  what is the quantitative requirement,

ii. what is the scientific justification, and

iii. what are the consequences of not meeting the requirement. [^1]

In each of the following requirements we address these questions.

At the end of the chapter we will list success criteria: complete,
substantial and minimal. The quantitative requirements are those that if
met allow complete success, whereas looser requirements that allow for
substantial success may be suggested in the consequences section.

## Survey area and survey time span

1.  The DES will cover 5000 square degrees of sky in the South Galactic
    Cap, of which 4000 sq. degrees will overlap the SPT survey area and
    the remainder will be devoted to maximizing overlap of DES
    photometry with redshift surveys. The SN survey will cover 40
    sq-degrees observed using 10% of the available time.

2.  The survey should be completed in 5 years using 30% of the available
    telescope time, taking account of weather.

3.  The area should have the minimum practical perimeter/area ratio.

4.  The area should be at reddening 0.2 mag of *r*-band extinction and
    stellar density ≤ 10,000 stars/sq-degree.

*Science justification*

A large area of homogeneous photometry is required to produce the
extensive galaxy map that forms the core data set for cluster finding,
weak lensing, and angular power spectra of galaxies (S-1). Evaluations
of expected constraints suggest that the return on the investment for
cluster surveys is maximized near a fractional sky coverage of 0.1 or
4000 sq-degrees (S-1).

Overlapping the SPT survey area will enable the use of SZ as cluster
locator and mass estimator, and will enable the high quality science of
the SPT survey.

This amount of exposure time for the supernova survey will be sufficient
to detect and measure the target goal of 2000 objects (S-1 & S-2).

We are constrained in time by the NOAO allocation. A survey designed to
need more time is unlikely to be accepted by the wider astronomical
community, although additional time may be applied for via the standard
NOAO procedures.

Regions near the edges of the survey are harder to calibrate (S-3). A
larger ratio of edge to survey area effectively reduces the coverage.

Greater reddening limits the depth constant exposure time achieves, and
higher density of stars reduces the fraction of the image searchable for
galaxies; both of which hamper the construction of complete galaxy
catalogs to a specified magnitude (S-4).

As the DES is not a spectroscopic redshift survey, we rely on other data
sets to provide the photo-z calibrations. To be useful, the survey must
cover this area with high quality photometry.

*Consequences of not meeting requirements*

The constraints on w depend on the square root of the surveyed area. Not
surveying the entire survey area will decrease the precision of our
measurements; to a small extent for small shortfalls, to a much greater
extent for shortfalls of order 50%.

Insufficient area covering photo-z will make calibration difficult,
lessen our ability to understand the photo-z error distribution and thus
lessen our ability to interpret the weak lensing results.

## Wavelength coverage

5.  The DES will cover the wavelength range from 400nm to 1000nm using 4
    filters.

*Science justification*

To meet the photometric redshift requirement we will bracket the 4000Å
break signal in the elliptical galaxy spectrum typical of cluster
galaxies. At z = 0, the break is at 390 nm while at z = 1, the break is
at 800nm; redshift moves the 4000Å break to the red. Precision
photometric redshifts require a pair of filters, one containing the
break and one redward of the break. We thus need bandpasses from 400nm
to 1000nm. We adopt an SDSS-like filter set, *g, r, i, z*. This choice
allows us to exploit the existing body of work done with these filters
to aid in analysis and calibration of DES data.

We do not require exact matching of the SDSS system and will calibrate
off the spectroscopic redshifts obtained from a region of overlap with
existing redshift surveys.

*Consequences of not meeting requirements*

Photometric redshifts are possible with three filters, though using
three instead of four would lessen the precision of the photo-z,
increase the catastrophic error rate, and lessen our ability to minimize
the systematic errors. These translate into lower precision on w, with
some of the more sophisticated weak lensing analyses being the most
sensitive.

## Limiting magnitudes

6.  The DES will reach to limiting AB magnitudes for galaxies of *g* =
    24.0, *r* = 24.0, *i* = 24.0, and *z* = 23.6.

*Science justification*

Limiting magnitudes are driven by two science goals: detecting
(reliably) 1/2 L\* galaxies in galaxy clusters out to redshift z = 1,
and obtaining photometric redshifts of galaxies to a density of 20 per
sq-arcmin.

The DES will, for all redshifts at 0.1 ≤ z ≤ 1, obtain 10ff magnitudes
in the filter bandpass containing the redshifted 4000Å break, and in the
adjacent redward filter, for a 1/2L\* galaxy with a passively evolving E
spectrum. A limiting absolute magnitude corresponding to a 1/2L\* galaxy
suffices to obtain photometric redshifts of clusters with high precision
(σ~z~ = 0.01) and also enables the construction of high completeness and
purity cluster catalogs with minimal selection functions. The absolute
magnitude of an L\* galaxy at z = 0 is M~i~ = -21.0 in the i-band
(Blanton et al 2003, redshift and evolution corrected). At z = 1 such a
galaxy would have a *z*-band magnitude of 23.3 assuming the galaxy
exhibits the passively evolving old stellar population model appropriate
for high luminosity cluster galaxies.

The DES will go 0.3 magnitudes deeper in z to ensure that the galaxy
catalog is complete at 23.3.

The DES will obtain sufficiently deep *g, r* data such that for every
object detected in *z* band with a flat or bluer spectrum the
photometric redshift will have σ~z~ ≤ 0.2 (barring catastrophic error).

The minimum apparent magnitudes in the remaining bandpasses follow by
considering the redshifted, passively evolving spectrum of a cluster
galaxy. The limiting magnitude required is that for an L\* galaxy at the
redshift where the 4000Å break leaves the bluer filter and enters the
redder filter. The 4000ºA break leaves the *g, r* and *i* bandpasses at
redshifts of 0.35, 0.65, and 1.0 respectively, and there we find
\[*g,r,i*\] = \[22.8, 23.4, 24.0\].

Data from both the CTIO Blanco 4-m and the CFHT 3.6m show that 10*σ* is
near where the differential number counts turn over, an indication that
it is close to the completeness limit. We will go 0.3 mag deeper in the
fiducial band in order to safely construct a flux-limited sample of
galaxies. We therefore require a limiting magnitude of *z* = 23.6.

The minimum apparent magnitudes required for decent photometric
redshifts for blue galaxies is fainter than that required for the red
galaxies.

Table 2.1 Progression of Required 10*σ* Limiting Magnitudes

  ---------------------------------- --------- --------- --------- ---------
  Filter                             g         r         i         z

  clusters                           22.8      23.4      24.0      23.3

  Galaxy completeness                22.8      23.4      24.0      23.6

  Blue galaxy photo-z                24.0      24.0      24.0      23.6
  ---------------------------------- --------- --------- --------- ---------

The depth of images is determined by the aperture of the telescope, the
exposure time, the transmission of the atmosphere, the sky brightness,
the reflectivity or transmission of all relevant optics, the quantum
efficiency and noise properties of the CCD detectors, the seeing and
image quality, and the efficiency of the reduction software.

The technical specifications contain an estimate of these quantities
(including a CCD readnoise of less than 15 e-), we find that translating
the limiting galaxy magnitudes into 5*σ* detection limits for a point
source in 0.9\" seeing gives *g* = 25.9, *r* = 25.4, *i* = 25.5, and *z*
= 24.9.

The 5σ point source limit can be made into a requirement on throughput
and quantum efficiency: A 21.0 magnitude star imaged for 100 seconds at
an airmass of 1.3 under photometric conditions should give more than
20,000, 20,000, 15,000, and 15,000 e-, for the four filters,
respectively.

*Consequences of not meeting requirements*

Not meeting the requirements will make the construction of the cluster
catalog more difficult, and the interpretation of the results harder.
The ultimate precision on w will suffer.

The weak lensing survey will be less effective, as the effective number
of galaxies per sq-arcminute is a fairly strong function of limiting
magnitude. The ultimate precision on w will suffer.

The ability to track the light curves of high-z supernovae is crucial to
extracting constraints on w in this method. Given the cadence demands of
the supernova techniques, the depth of the survey in relatively short
exposures is crucial.

## Point spread function

7.  The instrument plus telescope should deliver images no worse than
    0.55'' for zenith-corrected images in the *r, i*, and *z* bands.

8.  

9.  The weak lensing galaxy shear measurements, averaged over 2 degree
    field of view, should be statistics limited, with systematic
    calibration error increasing the error by no more than 10%.

10. During normal operations, the systematic error in the calibration of
    the point spread function (PSF) width and ellipticity shall be less
    than 0.1% averaged over scales of 2 degrees.

*Science justification*

The science goals of the DES are controlled in part by the image
quality, including the number of objects detected, the efficiency of
star-galaxy separation at all magnitudes, and the efficiency of using
galaxies for weak lensing. Good seeing clearly pushes all these
requirements in the same direction. The survey should be limited by site
seeing, not instrument performance, during typical observing (S-7, S-8).

The DES strategy has been derived assuming that the statistics on image
sizes from the DES camera are the same as those attained today with the
existing Mosaic II imager and Prime Focus corrector. Only images with
FWHM ≤ 1.1'' (≤ 1.2\" in *g*) will be used in the coadds. The survey
strategy presumes that 90% of all data collected on photometric nights
will meet the imaging PSF requirement.

The current total contribution of the telescope and Mosaic II imager to
the PSF is estimated to be about 0.55" FWHM. This value has been
determined by comparing long-term measurements of the site seeing from
the RoboDIMM monitor with long-term statistics on the measured FWHM from
the Super-Macho survey.

The science reach of the DES is increased as the image size is made
smaller up to the ultimate limit set by the atmosphere, the figure of
the primary mirror, and lateral charge diffusion in the CCDs. The
enhanced goal is intended to indicate what would be a practical limit
for what would constitute the "best\" images.

The weak lensing science goal drives the requirements as the effects of
large scale structure on shear measurements is at the 1% level. The
density of galaxies useful for weak lensing measurements is a strong
function of PSF size.

Stars in the science images are used to measure the PSF. In a single 100
sec exposure we can reach magnitude 21 at S/N ≈ 100. At this magnitude
there are about 2000 stars per sq-degree. Based on experience with SDSS,
roughly a quarter of these 2000 stars will be away from frame edges,
cosmic rays, and other stars and galaxies. This corresponds to a little
more than 1 useful star per sq-cm (= 9 sq-arcmin). With this
information, a map of the PSF and possible variations over the focal
plane will be generated. One can sample the PSF on scales smaller then 1
sq-cm by averaging over several images but only if the instrumental PSF
is constant against, for example, gravity induced flexure changes or
thermal effects (S-10).

There are approximately 100,000 galaxies in a single 3 square degree FOV
image (10 galaxies/sq-arcminute to the limiting magnitude). The shear
signal on the scale of 2 degrees is approximately 1%. The statistical
error in measuring this signal (averaging over shape noise) is expected
to be 0.1%. For galaxies 1.2" in size, the statistical error in
"whisker" length is 0.06 arcsec. (A whisker is: (*a*^2^ - *b*^2^)^1/2^
where *a*^2^ and *b*^2^ are the major and minor axis second moments of
an image.) The PSF calibrated from stars averaged over the field must
have a peak statistical error of no more than 0.06" in equivalent
whisker length (S-10).

The decision to not include an ADC is driven in part by science
considerations. Weak lensing will be done primarily in the *i* and *z*
bands, where differential refraction is minimal. An ADC introduces
zenith angle-dependent systematic effects that mimic a weak lensing
signal and are comparable in magnitude to differential refraction, and
thus the ADC effects would need to be calibrated on top of all other
effects. The strategy in existing CTIO and KPNO weak lensing surveys has
been to keep the ADC turned off. We would do so the same. The tradeoff
is that images obtained in good seeing and at high zenith angles,
particularly in the *g* band, would be somewhat elongated. We estimate
that the elongation will be 10% (for 1" seeing) in the *g* band at the
highest airmass used in the survey (sec z = 1.5). No science goals would
be threatened ().

*Consequences of not meeting requirements*

Weak lensing measurements depend critically on being able to measure the
shape of the PSF, and the stability of the PSF so as to carry the shape
measurement in space and time. The number of galaxies useful in the
measurements depends on the size of the PSF relative to the size of the
galaxies as well as the S/N on the shape measurements. If the PSF is
varying rapidly across the CCDs, then the effective noise in the
measurement goes up as the uncorrected PSF shapes contribute noise.

Likewise, star-galaxy separation and the depth of the data will all be
strongly affected.

Lastly, the limiting magnitude of an image is a strong function of the
seeing.

## Throughput

11. The average throughput of the telescope and corrector will be \>=
    60%, averaged over the focal plane.

12. The photon detection efficency of the filters and CCDs will be \>=
    40% for g and \>= 55% for r,i,z, averaged over the entire bandpass
    and the entire CCD.

*Science Justification*

These two throughputs were assumed in the calculations that support the
feasibility of the required limiting magnitudes (S-6) over the required
survey area (S-1) given the availability of the telescope (S-2).

The telescope is assumed to have an effective area of 10 sq-meters. The
throughput on the telescope and corrector then is a cumulative effect of
the reflection off aluminum, transmission through the glass of the
corrector elements, and the effects of the anti-reflection coatings on
the corrector elements.

*Consequences of not meeting the requirements*

These two requirements may be traded off- it is the system throughput
that matters when trying to achieve the required limiting magnitudes.

If the combined throughputs are less then required then we have two
options, either surveying less area or achieving a shallower limiting
magnitude.

## Photometric redshifts

13. The DES will obtain σ~z~ ≤ 0.01 for groups and clusters of galaxies
    at 0.1 ≤ z ≤ 1.0

14. The DES will obtain σ~z~ ≤ 0.03 for red galaxies at z ≤ 1

15. The DES will obtain σ~z~ ≤ 0.2 for blue galaxies with flat or bluer
    spectra detected in *z*-band, and z ≤ 1.

16. The rms of errors in photometric redshifts and the residual bias in
    the photometric redshifts will be measured.

*Science justification*

All of the science goals require precision photometric redshifts.

To achieve the required understanding of the systematic uncertainties a
calibration sample of roughly 1000 redshifts per 0.1 bin in redshift is
needed in addition to the DES data. Our strategy is to make use of
public spectroscopic datasets in areas of the sky that we can image. The
largest and most diverse of these datasets is that of the SDSS Southern
Survey which has 100,000 redshifts down to *i* = 20 and out to z = 0.5,
in combination with the 2dF-SDSS survey, which has 10,000 redshifts of
red galaxies out to z = 0.75. This same area of sky also contains fields
from both the VIRMOS VLT Deep Survey (8 sq-deg with spectroscopic
redshifts to magnitude I~AB~ = 22.5, and 2 sq-deg with redshifts to
I~AB~ = 24) and the Keck Deep2 Survey (2 sq-deg with spectroscopic
redshifts to R~AB~=24). These surveys will deliver tens of thousands of
spectroscopic redshifts extending over our redshift range of interest.

*Consequences of not meeting requirements*

We can suffer a factor of 2 increase in the errors as long as they
remain substantially Gaussian, without too much decrease in our
sensitivity to w. Many techniques will bin the galaxies into 5 or 10
shells between 0 ≤ z ≤ 1, and will not gain any advantage from thinner
shells. A few of the techniques based on correlations can benefit from
higher precision photometric redshifts, but these tend to be the ones
most sensitive to the details of the error distribution. Finally, the
very high precision cluster photometric redshifts that seem possible, of
order σ~z~ ≈ 0.003 may well open up a new realm of analyses.

Most of the worry about lowering the precision of the photo-z
requirements is based on the worry that the error distribution will
become increasingly non-Gaussian.

## Photometric calibrations

17. The magnitudes of an object may be calculated to within 2% by
    convolving the spectrum of the object with the system response
    curves. This requirement assumes that the spectra are
    spectrophotometrically calibrated and that the system response
    curves are absolute.

18. The magnitudes vary only by --2.5 log *f*~2~/*f*~1~, independent of
    position in the final map to within 2% (1% enhanced goal), where
    *f*~2~/*f*~1~ is the ratio of photon fluxes. This is to be true in
    *g, r, i, z* individually.

19. The magnitudes have an absolute zero point that is well-defined and
    known to 0.5%. The magnitudes will be on the natural instrument
    system.

*Science justification*

The system response curve requirement (S-17) is driven by several
considerations:

a)  one of the two major branches of photometric redshift technologies
    uses template spectra convolved with system response curves;

b)  the interpretation of the division of the galaxy catalog into color
    and redshift bins requires that we know how to predict the colors;
    and

c)  supernova observations through a filter must be corrected to a rest
    frame magnitude, which requires very good knowledge of the system
    response.

It is not sufficient to use a mean CCD and filter transmission curves as
it is known that both CCDs and filters can a) show significant
differences, and b) change with time. We are not trying to match the
SDSS calibration.

Both the relative and absolute calibration requirements (S-18, & S-19)
are driven by the amplitude of the large scale structure that the survey
will measure. Errors in photometric zero point cause errors in galaxy
counts that mimic large scale structure in any apparent magnitude
limited sample. On angular scales comparable to that of the whole
survey, an error budget to the galaxy number counts is:

-   2% true large-scale structure, to be compared to

-   4% Galactic extinction if uncorrected, and

-   2% systematic photometric error.

Current generation extinction maps can be used to correct for Galactic
extinction. The photometric calibration is required to be 2% in each
bandpass based on photometric redshift requirements, though as an
enhanced goal, we expect to calibrate the final coadded map to 1%.
Further, we require 3% calibration at the important two year milestone.
The accuracy of the calibration can be verified by checking the position
and scatter in the stellar locus.

We have no science requirement on the calibration of single images, so
we choose a number that makes them useful for the wider community. We
estimate that these can be calibrated to 10% using the USNO B2 star
catalogue stars that will be in each frame.

The 2% (or 0.02 mag) allowed error in the magnitudes as defined above
has several contributions; the error from each contribution is
arbitrarily taken to be 2% **∕** √2, although ultimately it is the
combination of all that must be constrained. For a few of these
contributions, we use a smaller error, where it is justified.

### Linearity: 

20. Uncorrected nonlinearities due to imperfect shutter timing and
    nonlinear CCD/amplifier gain shall be less than 0.3%, measured as
    the peak error between shortest and longest exposure times, and
    between the faintest and brightest unsaturated stars.

The standard stars are observed with short integrations at high count
rate, while the survey data are observed with longer integrations at low
count rate. Any nonlinearities will give systematic errors, an offset in
the absolute calibration and a scaling problem in the relative
photometry.

### Ghost subtraction: 

21. The ghost image of the night sky on the CCDs will be removed to no
    worse than 0.6%.

Ghosts input structure on the flat field that will not subtract using
local sky. They may be removed if well characterized.

### Flatfielding: 

22. The rms variation in photometric calibration of any star measured in
    different positions of a CCD shall be no worse than 0.6%.

Standard stars are observed in a small area of the CCD, while secondary
stars cover the entire area.

### Residual stray light subtraction: 

23. The remaining stray light in the system will be no worse than 0.6%.

Stray light input structure on the flat field that will not subtract
using local sky, structure that affects the photometry. Stray light
subtracts out in local sky subtraction. Flat fields do not have sky
subtraction: scattered light puts false structure into the flat fields.
Ghosts are extreme scattered light, but they may be modeled and removed.
Both some and sky flats are affected, but scattered light from stars are
an issue in the sky flats.

### Variations in aperture correction: 

24. The aperture correction shall have an internal rms error no bigger
    than 0.6% for any CCD and seeing between 0.8" and 1.5".

The aperture correction converts PSF counts to the counts inside a large
aperture. The PSF will vary in any image (and coadded image), so this
quantity will vary across a field. The photometric pipeline will take
this variation into account as best it can. We may be able to do
substantially better than this requirement.

### Primary standard star calibration: 

25. The primary standard star set we choose as the absolute calibration
    sample should have *g, r, i, z* magnitudes known to \< 0.5%.

The absolute calibration should not dominate the calibration error
budget.

*Consequences of not meeting requirements*

Systematic magnitude errors will strongly affect the angular power
spectra key project, and cluster finding and weak lensing to strong will
be affected in serious ways. Random magnitude errors, if they are
quantified, can be taken into account in large-scale structure studies,
but gradient pose great problems.

We can tolerate 3% photometry with minor loss of precision in the
photometric redshifts. Photometry worse than 3% will seriously limit all
of the key science goals of the survey.

## Astrometric calibration

26. The absolute astrometry on the sky should be no worse than 100 mas.

27. The relative astrometry between adjacent colors should be no worse
    than 100 mas.

*Science justification*

The survey must average together images to reach the necessary depths.
This averaging process should not degrade the PSF; in fact, it should
help it but only of the images of the objects are aligned to better than
a third of a pixel, 0.1\".

Relative astrometry of 100 mas is ideal for deblending, and also
minimizes systematic errors in aperture photometry colors.

Available all-sky/wide-area astrometric catalogs can be used with DES
data to determine proper motions. The 100 mas requirement ensures that
the accuracy of the proper motion determinations will be limited by the
other catalogs, not by DES. Additionally, DES data will be used to
identify optical counterparts of objects detected in other spectral
regions. Again, this requirement will ensure that DES positions are not
the limiting factor in making the identifications.

The requirements apply to the positions determined in *i* for
point-source objects brighter than *i* = 22 with spectral type between
O5 and M0. The determination of positions for very blue, very red, or
emission-line objects may, and certainly occasionally will, fail to meet
the requirements, due to differential chromatic refraction.

Note that the CCD is an optical element, with distortions possible on
the 0.3 pixel scale near substrate vias.

*Consequences of not meeting requirements*

If the astrometry is worse than 0.1\" the coaddition process will start
degrading the PSF, introducing first structure into the shape of the PSF
and then increasing the size of the PSF. These effects will be on the
spatial scale of the errors in the astrometry. This will have strong
consequences for the weak lensing science.

The deblender, as well, will have to be run with much looser constraints
on distinguishing close blends, with strong consequences for studies of
close pairs of objects, and rejection of moving objects

We could probably tolerate 0.2\" astrometry at the loss of precision in
the weak lensing measurements. Any worse than this would be unacceptable
for both lensing and galaxy photometry.

## Mapping

28. The input images are to be mapped into identical locally flat images
    ("charts\") for coaddition, where the distortion free point
    ("tangent point\") is at chart center.

29. The distortion in the charts is to be less than 0.1%, peak to peak.

30. The mapping should conserve flux.

31. Mask images accounting for bad data (e.g., bad columns, bleed
    trails, saturation) will also be mapped.

*Science justification*

To make the terminology concrete: the input images are mapped to charts,
and the charts are coadded into a final map. The problem is that there
is no ideal projection of the sphere onto a plane, where ideal would
preserve lengths, angles and areas. In the end, this is due to the fact
that spherical triangles sum to \> π radians and planar triangles sum to
π radians. It is possible to create projections that preserve angles
("conformal projections\") or that preserve area ("equal area
projections\"); the latter are more likely to be useful to us.

The charts are created so as to make coaddition well defined. Since the
coaddition should not introduce artifacts the charts should, for any
coadd center, be free of distortion over the tolerance required. This is
0.1\" for the science maps. In the creation of the charts from the data,
flux should not be lost or arbitrarily scaled. In other words, the
charts should preserve the photometry of the input images. This issue
has three components:

i.  Flux conservation during the remapping. This requires that the
    Jacobean matrix, describing the change of coordinates, is known. It
    should be reversible.

ii. Flux conservation during the mapped pixel estimation. This requires
    that the convolution kernel used to estimate the chart pixel is
    normalized.

iii. Total flux conservation. This requires that every pixel in the
     input image that transforms to the interior of the chart is sampled
     at least once.

Masks will be important to dismiss bad data from the coadd, optimizing
the result.

*Consequences of not meeting requirements*

If the mapping is worse than 0.1\" the coaddition process will start
degrading the PSF, introducing first structure into the shape of the PSF
and then increasing the size of the PSF. These effects will be on the
spatial scale of the errors in the distortion of the mapping. This will
have strong consequences for the weak lensing science.

The deblender, as well, will have to be run with much looser constraints
on distinguishing close blends, with strong consequences for studies of
close pairs of objects, and rejection of moving objects

We could probably tolerate 0.2\" astrometry at the loss of precision in
the weak lensing measurements. Any worse than this would be unacceptable
for both lensing and galaxy photometry.

Flux conservation problems translate directly into photometric
calibration problems. There should be no need to lose flux at anything
larger than the round-off errors, but a practical limit would be 0.1%.

##  Summary of Science Requirements

The full list of science requirements is as follows.

1.  The DES will cover 5000 square degrees of sky in the South Galactic
    Cap, of which 4000 sq. degrees will overlap the SPT survey area and
    the remainder will be devoted to maximizing overlap of DES
    photometry with redshift surveys. The SN survey will cover 40
    sq-degrees observed using 10% of the available time.

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32. The survey should be completed in 5 years using 30% of the available
    telescope time, taking account of weather.

33. The area should have the minimum practical perimeter/area ratio.

34. The area should be at reddening 0.2 mag of *r*-band extinction and
    stellar density ≤ 10,000 stars/sq-degree.

35. The DES will cover the wavelength range from 400nm to 1000nm using 4
    filters.

36. The DES will reach to limiting AB magnitudes for galaxies of *g* =
    24.0, *r* = 24.0, *i* = 24.0, and *z* = 23.6.

37. The instrument plus telescope should deliver images no worse than
    0.55'' for zenith-corrected images in the *r, i*, and *z* bands.

38. The weak lensing galaxy shear measurements, averaged over 2 degree
    field of view, should be statistics limited, with systematic
    calibration error increasing the error by no more than 10%.

39. During normal operations, the systematic error in the calibration of
    the point spread function (PSF) width and ellipticity shall be less
    than 0.1% averaged over scales of 2 degrees.

40. The average throughput of the telescope and corrector will be \>=
    60%, averaged over the focal plane.

41. The photon detection efficency of the filters and CCDs will be \>=
    40% for g and \>= 55% for r,i,z, averaged over the entire bandpass
    and the entire CCD.

42. The DES will obtain σ~z~ ≤ 0.01 for groups and clusters of galaxies
    at 0.1 ≤ z ≤ 1.0

43. The DES will obtain σ~z~ ≤ 0.03 for red galaxies at z ≤ 1

44. The DES will obtain σ~z~ ≤ 0.2 for blue galaxies with flat or bluer
    spectra detected in *z*-band, and z ≤ 1.

45. The rms of errors in photometric redshifts and the residual bias in
    the photometric redshifts will be measured.

46. The magnitudes of an object may be calculated to within 2% by
    convolving the spectrum of the object with the system response
    curves. This requirement assumes that the spectra are
    spectrophotometrically calibrated and that the system response
    curves are absolute.

47. The magnitudes vary only by --2.5 log *f*~2~/*f*~1~, independent of
    position in the final map to within 2% (1% enhanced goal), where
    *f*~2~/*f*~1~ is the ratio of photon fluxes. This is to be true in
    *g, r, i, z* individually.

48. The magnitudes have an absolute zero point that is well-defined and
    known to 0.5%. The magnitudes will be on the natural instrument
    system.

49. Uncorrected nonlinearities due to imperfect shutter timing and
    nonlinear CCD/amplifier gain shall be less than 0.3%, measured as
    the peak error between shortest and longest exposure times, and
    between the faintest and brightest unsaturated stars.

50. The ghost image of the night sky on the CCDs will be removed to no
    worse than 0.6%.

51. The rms variation in photometric calibration of any star measured in
    different positions of a CCD shall be no worse than 0.6%.

52. The remaining stray light in the system will be no worse than 0.6%.

53. The aperture correction shall have an internal rms error no bigger
    than 0.6% for any CCD and seeing between 0.8" and 1.5".

54. The primary standard star set we choose as the absolute calibration
    sample should have *g, r, i, z* magnitudes known to \< 0.5%.

55. The absolute astrometry on the sky should be no worse than 100 mas.

56. The relative astrometry between adjacent colors should be no worse
    than 100 mas.

57. The input images are to be mapped into identical locally flat images
    ("charts\") for coaddition, where the distortion free point
    ("tangent point\") is at chart center.

58. The distortion in the charts is to be less than 0.1%, peak to peak.

59. The mapping should conserve flux.

60. Mask images accounting for bad data (e.g., bad columns, bleed
    trails, saturation) will also be mapped.

These science requirements may be summarized as follows:

i.  5000 sq-degrees in the South Galactic Cap, significantly overlapping
    the SPT SZ survey, to be completed in 5 years with 30% duty cycle.

ii. Four bandpasses covering 390nm to 1000 nm: *g, r, i, z.*

iii. Limiting magnitudes of \[*g, r, i, z*\] = \[24, 24, 24, 23.6\] for
     small galaxies.

iv. A PSF of FWHM ≤ 1.1\"and median ≤ 0:9\", stable over 10 sq-arcminute
    areas.

v.  Photometric redshifts to z = 1 for clusters, red galaxies, and blue
    galaxies with precisions of σ~z~ = 0.01; 0.03; 0.20 and understood
    error distributions

vi. Photometric calibration to 2% with understood system response curves

vii. Astrometric calibration to 0.1".

viii. Image remapping that conserves flux and is practically distortion
      free.

ix. A galaxy catalog containing measurements that allow straightforward
    selection function calculation and simulation.

x.  Collaboration scientists have high availability and easy access to
    survey data.

##  Success criteria

We define two levels of success.

*Complete success*

Complete success is defined as meeting all of the criteria of section
2.10 or, in\
summary, as follows

5000 sq-degrees in the South Galactic Cap, significantly overlapping the
SPT SZ survey, to be completed in 5 years with 30% duty cycle.

i.  Four bandpasses covering 400nm to 1000 nm: *g, r, i, z*.

ii. Limiting magnitudes of \[*g, r, i, z*\] = \[24, 24, 24, 23.6\] for
    small galaxies.

iii. A PSF of FWHM ≤ 1.1\"and median ≤ 0:9\", stable over 10
     sq-arcminute areas.

iv. Photometric redshifts to z=1 for clusters, red galaxies, and blue
    galaxies with precisions of σ~z~ = 0.01; 0.03; 0.20 and understood
    error distributions

v.  Photometric calibration to 2% with understood system response curves

vi. Astrometric calibration to 0.1\".

vii. Image remapping that conserves flux and is practically distortion
     free.

viii. A galaxy catalog containing measurements that allow
      straightforward selection function calculation and simulation.

ix. Collaboration scientists have high availability and easy access to
    survey data.

x.  Collaboration analysis papers on all key projects.

*Substantial success*

Substantial success is defined to be a survey that meets at least the
following requirements:

i.  4000 sq-degrees in the South Galactic Cap, significantly overlapping
    the SPT SZ survey, to be completed in 5 years with 30% duty cycle.

ii. Three of four bandpasses covering 400nm to 1000 nm: *g, r, i, z*.

iii. Limiting magnitudes of \[*g, r, i, z*\] = \[23.5, 23.5, 23.5,
     23.3\] for small galaxies.

iv. A PSF of FWHM ≤ 1.3\"and median ≤ 1.0\".

v.  Photometric redshifts to z=1 for clusters, red galaxies, and blue
    galaxies with precisions of σ~z~ = 0.03; 0.06; 0.40 and understood
    error distributions

vi. Photometric calibration to 3% with understood system response curves

vii. Astrometric calibration to 0.2''.

viii. Image remapping that conserves flux and is practically distortion
      free.

ix. A galaxy catalog containing measurements that allow straightforward
    selection function calculation and simulation.

x.  Collaboration scientists have high availability and easy access to
    survey data.

xi. Collaboration analysis papers on all key projects.

# Technical Specifications

To meet the science requirements listed above we must make choices that
strike a balance between the science requirements and what is
technically feasible. Below we present our specification for the CCDs,
the integration time of the images, readout time and read noise, the
pixel size, image quality required for the optics, and the size of the
field of view of the corrector and camera. These lead to a set of
technical requirements that are used to guide the design of the DES
corrector and camera.

## Survey Footprint

The DES will cover a total of 5000 sq. degrees in the Southern Galactic
Cap. Figure 1 and Table 3.2 show the specific areas. Below we describe
the motivation for these choices.

Table 3.2 Dark Energy Survey areas

  ----------------- ----------------- ----------------- -----------------
  Overlap target    Right Ascension   Declination Area  

                    (deg)             (deg)             (sq. deg.)

  SPT               -60 to 105        -30 to -65        3988

                    -75 to -60        -45 to -65        

  SDSS Stripe 82    -50 to 50         -1 to 1           200

  Connection region 20 to 50          -30 to -1.0       800
  ----------------- ----------------- ----------------- -----------------

*Justification*

1\. The SPT Survey Area is predefined

2\. SDSS Stripe 82 has been identified as the optimal area for DECam
photometric redshift calibrations via the SDSS spectroscopic redshifts.

## Field of view

The size of the field necessary to meet the science requirements depends
on many aspects of the survey. The primary constraint is to cover the
5000 sq-deg DES area in the time available for the DES during the 5 year
period. However, consideration must also be given to factors of weather,
survey strategy, and operational efficiency, among others.

![](media/image1.png){width="3.5570866141732282in"
height="3.4661417322834644in"}

Figure 1: The DES Footprint. The image shows a Lambert\'s equal area
projection of the DIRBE dust map covering the South Galactic cap. The
graticule shows celestial coordinates, with the South Celestial Pole
near the bottom. The DES footprint is in red.

An order of magnitude calculation shows the essentials. We are proposing
to use 30% of the Blanco 4-m time over 5 years, resulting in a survey of
525 nights. Our supernova program will take 10% of the time, leaving 473
nights for the imaging survey of 5000 sq-degrees. This requires an
acquisition rate of 10.6 sq-degrees per night. The nights in the
Southern summer will be around 8 hours long, giving a 10.6 sq-degree/8
hour rate. Assuming half the nights will be unusable due to bad weather,
this calculation gives 10.6-sq-degrees/4 hours. Our total exposure time
on a given field will be of order an hour, so the camera must cover 2.7
sq-degrees.

A more careful analysis (see Chapter 4 of the NOAO proposal),
incorporating overheads, a four-bandpass calibration strategy, and a
detailed weather model, results in the need for a camera whose single
image tiles cover a footprint of 3.0 sq-degrees and have a 10% dead area
for inter-CCD gaps. This footprint requires that the circumscribed
circle of the field of view have diameter of 2.2 degrees. The science
and operational requirements that drive this analysis are as follows.

1.  Number of tilings - The minimum number of tilings of the survey area
    in any filter is 4.

Multiple tilings are needed to fill in the inter-CCD gaps and to achieve
the required photometric precision. This number of tilings leads to the
*g* and *r* bands having deeper limiting magnitudes than is need by the
science requirements.

2.  Integration Times -- \[*g, r, i, z*\] = \[400, 400, 1200, 2000\]
    sec.

3.  Single exposure Time -- minimum 100 sec

To achieve our science goals we require a signal-to-noise of 10 at the
limiting magnitudes. The integration time is the total of all individual
exposures. The calculation of the exposure time depends on the mirror
area, the throughput of the system, the read noise of the instrument,
the sky background, and the area over which the object is spread.

The sky brightness contributes to the noise in galaxy magnitudes and
colors in proportion to the area of the aperture one uses to measure the
magnitudes. We assume the seeing at the Blanco 4-m is 1.0\", and use an
aperture of 1.7\" diameter, as is appropriate for the measurement of the
colors and magnitudes of faint galaxies. The resulting minimum
integration times are listed in Table 3.3.

Table 3.3 Limiting Magnitudes to Integration Times

  ----------------------------------------- ------ ------ -------- -------
  Filter                                    *g*    *r*    *i*      *z*

  Required limiting magnitude (10*σ*)       22.8   23.4   24.0     23.6

  CCD QE                                    0.65   0.85   0.65     0.5

  Filter and optics T Δλ/λ                  0.18   0.14   0.13     0.18

  Sky brightness (mag/sq-arcsec)            21.7   20.7   20.0     18.7

  Sky Brightness 45° from moon              19.7   19.7   19.25    18.5

  Minimum integration times (sec)           27     130    900      1600

  Adopted integration times (sec)           400    400    1200     2000

  Achieved Limiting magnitude (10 *σ*)      24.6   24.1   24.3     23.9
  ----------------------------------------- ------ ------ -------- -------

The last row gives our adopted integration times for the reference
survey strategy; please see Chapter 4 for a detailed discussion. The
increased exposure times in the *g* and *r* filters reflect our
photometric calibration strategy which will use multiple survey tilings
and a minimum exposure time of 100 sec.

4.  Observing time contingency: 25%.

The baseline observing plan should include a contingency of 25% (1
year). Although the plan factors in weather and seeing patterns at CTIO
and overhead for slew times between fields, it otherwise presumes that
operations are 100% efficient, thus excluding such factors as observer
errors, equipment malfunctions, etc. It also excludes the possibility of
anomalously bad weather years (such as consecutive El Nino years).

## Pixel scale

5.  Pixel scale: 2.5 pixels per FWHM, met by 0.27"/pixel, 15μ

The only science requirement on pixel scale is that images be properly
sampled (2.5 pixels per FWHM) under conditions of good seeing. This
requirement is met by the natural scale at the Blanco prime focus
(0.27"/pixel for 15 micron pixels). Larger pixels (up to 0.3"/pixel)
might also be acceptable, since the best images measured by the Blanco
telescope Mosaic II imager are 0.6", and one could cover more sky with a
single CCD. However, because CCD diffusion is the second largest
contributor to image size, one might instead want to reduce the pixel
size. Since corrector designs become more challenging as one adds power
(in either direction), we have not placed any constraints on pixel scale
other than that it remain at the nominal natural scale of the telescope.

## Point spread function

6.  The delivered PSF from the telescope, instrument, and other factors
    exclusive of the site seeing shall be no greater than 0.55 arcsec.

The delivered PSF is a combination of atmospheric and dome-induced
seeing, CCD diffusion, the coorector optical design, mechanical
alignment and stability, as well as the telescope tracking (guiding and
focus). The requirement on image quality, exclusive of atmospheric
seeing, is that it be no worse than that delivered by the current
telescope and mosaic imager system.  The following table gives a
preliminary budget for contributions to the image quality, along with
the basis for each contribution.  CCD diffusion, the largest
contribution, is the value currently achieved by LBNL with CCDs matching
the DES CCD design.  The optical corrector itself is required to produce
images comparable to the existing as-designed corrector.

The combined contribution is 0.47 arcsec, smaller than the requirement
of 0.55 arcsec.  The remaining 0.29 arcsec is carried as a contingency
for remaining contributions that are not yet identified or are difficult
to quantify at present.

The delivered PSF is a combination of atmospheric and dome-induced
seeing, CCD diffusion, the optical design, mechanical alignment and
stability, as well as the telescope tracking (guiding and focus). The
primary mirror figure contributes about 0.16" and CCD diffusion
contributes about 0.37" FWHM (1-D sigma of 9μm for a scale of 17.7"/mm)
to the image size. The best quartile of seeing delivered by the site is
about 0.6". The combination produces images with a PSF width of 0.73".
We require that the corrector degrade this by no more than 10%, or about
0.33". \[This value corresponds to an rms radius of 11 - 12μm. It should
be noted that the current Blanco prime focus corrector design produces
images with an rms radius of 11μm.\]

Table 3.4 gives a preliminary budget for contributions to the image
quality.

The procedure for calibrating the PSF shape for weak lensing will likely
be developed as an iterative process. It is likely that data from
multiple images will be combined in order to increase the density of
stars sampling the focal plane. There are several possible sources of
systematic error that are comparable in amplitude to the target
systematic error of 0.06" whisker length. To make progress, the
following requirements will be made:

7.  No requirement is placed on the static PSF shape.

The variation in PSF shape across the focal plane is relatively
insensitive to details of the corrector designs that we are considering
and depends mainly on the PSF size itself. It will be presumed that the
combining of many images will be able to calibrate the variation in PSF
across the focal plane.

8.  For each dynamic source of systematic error (e.g., focus
    variations), the peak amplitude induced in the PSF whisker length
    will be no larger than 0.06".

It is assumed that the calibration process (e.g., using a PCA type
analysis), will partially calibrate each of these sources.

Table 3.4 Image quality budget

  ------------------ ---------- ----------- --------------------------------
  Source             FWHM\      RMS Radius  Reference
                     arcsec     (microns)   

  Dome Seeing        0          0           Not known with certainty

  Telescope Guiding  0.03       1           Guess - take same as focus
                                            errors

  Wind Shake         0          0           Assume \"calm\" night

  Corrector                                 

  Design             0.26       9.2         Current performance Blanco-2605

  Manufacturing      0.10       3.4         Radii, index, thickness,
                                            homogeneity, polishing, etc

  Silica             0.07       2.2         Grade C
  Inhomogeneity                             

  Assembly Errors    0.06       1.9         Decenter, tilt, etc.

  Flexure            0.04       1.2         Gravity loading, etc.

  Focal plan         0.05       1.6         30 micron p-p
  location                                  

  Thermal            0.05       1.6         -5 to +25 C, Steel

  CCD Diffusion      0.31       11          Assumes 7.5 microns rms 1-D,
                                            LBNL papers

  Depth-of-focus     0.03       1.0         Kubik and Estrada report (*i*
                                            band)

  Prim. mir. Figure  0.16       5.3         CTIO mirror testing report

  Prim. mir. support 0                      
  (static)                                  

  Prim. mir. support                        Assume small with adaptive
  (flexure)                                 optics

  Tel. collim.       0          0           Combine with flexure
  (static)                                  

  Tel. collim.       0.05       1.7         200 micron offset
  (flexure)                                 

  Focus              0.03       1           Scaled from SDSS 2.5 m focus
                                            loop performance

  **TOTAL**          **0.47**               **Telescope + Instrument**
  ------------------ ---------- ----------- --------------------------------

## Radioactivity and Charged Particle Events

9.  Charged particle events due to radioactive material in the
    instrument shall produce a background no worse than that arising
    from cosmic rays at sea level (0.01 events/cm^2^/sec).

The basis for this requirement is that, because the CCDs are thick
(250μm), charged particle tracks will deposit more charge and affect
more pixels compared with their impact on traditional, back-illuminated
devices. For a charged particle entering at an angle of 45 degrees, the
trail will be about 17 pixels long. Data processing software must
identify these trails as non-astronomical objects and remove them from
the image. Long exposures will need to be broken into shorter ones in
order to aid in distinguishing charged particle trails from astronomical
objects, decreasing the efficiency with which data are collected.

The main purpose of this requirement is to ensure that care is taken
with choice of materials to avoid inadvertent introduction of
radioactive materials, particularly in the vacuum dewar.

## Science CCD Specifications

Table 3.5 CCD dimensions

  --------------- -------------------------------------------------------
  Dimensions      33.316 x 63.390 mm

  Thickness       0.250 mm (tolerance is 30μm)

  Format          2048 x 4096 pixels, serial register is 2048 edge

  Pixel size      15μm×15μm
  --------------- -------------------------------------------------------

10. CCD nonlinearity: 1% or better

11. CCD full well: 130,000e-

12. CCD residual image: ≤ 10e- from 3×10^6^ e- spread over 5 pixels

13. CCD readout rate: 250 kpix/s

14. CCD dark current: ≤ 0.007e-/s/pix

15. CTE: \> 0.99999

16. CCD blooming control: not required

17. CCD prescan pixels: Not required, but acceptable

18. CCD frame transfer: not required

19. Outputs/CCD: 2

### Quantum efficiency

20. CCD quantum efficiency -- \[*g, r, i, z*\] : \[60%, 75%, 70%, 65%\]

The minimum required QE averaged over a device and bandpass is given in
Table 3.6. The maximum on-chip variation in QE (the flat field) should
be ≤ 5%. The maximum chip to chip variation in QE should be ≤ 10%.

Table 3.6 CCD QE Requirements

  ------------ -------------------------------- --------------------------
  filter       wavelength (nm)                  Required QE

  *g*          400-550                          60%

  *r*          560-710                          75%

  *i*          700-850                          70%

  *z*          830-1030                         65%
  ------------ -------------------------------- --------------------------

These came from the signal to noise calculations that support the 4 key
science projects. Filter transmission is also considered. See section
3.8.9.

### QE and Temperature Stability

A related issue is the QE stability. This is a requirement on the
temperature stability, as the temperature affects the red sensitivity of
the devices.

21. CCD QE stability: ≤ 0.3% variation in QE over 12-18 hours

Which implies:

22. CCD temperature stability - 1 degree K stability over 12-18 hours.

Which implies:

23. Spatial temperature variation across the focal plane should produce
    ≤ 5% variation in QE

Or:

24. Spatial temperature variation: ≤ 15 degrees K across the focal
    plane.

### Read noise

One of the components of the exposure time calculation is the read noise
of the devices. Because the collecting area of the Blanco 4-m is large,
the time required to reach the regime where sky noise dominates over
read nose is quite short. Read and sky noise add in quadrature. To meet
our science goals we require that the read noise be a factor of 2
smaller than the sky noise so that in the quadrature sum, the effect of
read noise is less than a 10% increase above pure sky noise.

25. CCD read noise must be ≤ 15e-/pixel

The *g* band has the darkest sky (thus smallest sky noise), and with our
100 sec exposures, a read noise of 15e-/pixel will be less than one-half
the sky noise. The CCDs we have chosen have been clocked at 250 kpix/sec
with a corresponding read noise of 7 e-. An additional contribution from
the front end electronics (system noise) will add in quadrature to this.
State of the art instrumentation can achieve \< 5 e- read noise from the
front end electronics. We choose 10 e- as our technical goal for the
total read noise. In practice, following modern design techniques, the
system noise will most likely be low enough to allow a slower scan
readout generating images limited by the lowest readout noise attainable
with these CCDs of 2-3 e-/pixel.

26. Goal for CCD read noise at DES readout rates: 10 e-/pixel

### CCD Diffusion and Flatness

27. The CCD charge diffusion should have an rms 1-D width *σ* no greater
    than 7.5μm.

This is the current performance of LBNL CCDs, and it is comparable with
the performance in other state-of-the-art CCDs (e.g., SITe, E2V).
Smaller widths can be achieve with higher substrate bias voltages but
would require additional CCD R&D effort.

The point spread function sets a requirement on the flatness of the
CCDs. The PSF should change by no more than 2% in second moment
semi-major axis and by no more than 2% in ellipticity, defined as b/a,
when convolved with a 0.6\" circular Gaussian and measured over 1 cm
spans. This translates to

28. CCD flatness: on \< 1 cm^2^ scales on a given CCD, the mean flatness
    variation should be ≤ 3μm.

29. CCD flatness: between adjacent ≤ 1cm^2^ areas on a given CCD, the
    difference in mean flatness should be ≤ 10μm.

This requirement derives from the ability of the weak lensing analysis
to track PSF variations; the last comes from the optical error budget.
Testing whether we meet the requirement involves mapping the height
variations and putting them into the optics model to evaluate the PSF
changes.

### Cosmetic defects

30. Cosmetic defects should account for no more than 5% loss in CCD area
    from non-usable pixels.

The requirement on bad columns derives from requiring no more than a 5%
loss in CCD area from non-usable pixels. This allows us to operate the
survey with 1 non-functional CCD in the focal plane. The budget is

-   1 non-operational focal plane CCD = 1.6%

-   10 cols ×2 lost due to glowing edge = 1%

-   10 rows ×2 lost due to glowing edge = 0.5%

-   each bad column affects entire 4096 column, and 5 columns centered
    on the bad one (PSF size is 2 pixels), so each bad column is a loss
    of 0.25%

-   1.9% allocated to bad columns translates to 8 bad columns, in the
    mean over the focal plane.

Adjacent bad columns count as 2 columns bad plus 4 more affected, 0.3%,
3 adjacent bad columns as 0.35% and so on. Hot pixels and traps are less
of a problem, if they don\'t affect many pixels.

## Telescope tracking, focus and active optics

The delivered PSF can be strongly affected by the telescope tracking and
focus. The TCS incorporates an interpolated lookup table model (pointing
model) to compensate for flexure and irregularities in open-loop
telescope tracking. Guide cameras follow stars near the object being
observed, sending tracking corrections to the TCS every second and
closing the tracking servo loop thus guaranteeing good quality images
for exposures of up to 1800 seconds in most observing programs.

31. Guide error signal update rate: 1 Hz.

Focus changes at the Blanco telescope resulting from changes in ambient
temperature are currently corrected manually by the observer. As noted
before, the existing prime focus camera delivers seeing typically
between 0.8\" and 1.1\". Unguided focus sequences, which are designed to
locate current best focus, show a median seeing of 0.7\" to 0.8\". It is
tempting to conclude that the telescope performs well on the 10 second
timescale of a focus exposure sequence, but less well in the several
hundred second exposures of a typical observing run. This would probably
be erroneous as many observers will happily and appropriately refocus
the telescope should the seeing improve, but not when it degrades when
it would be a waste of observing time. The existing guiding algorithms
are not perfect and effort will be put into refining them where
possible.

We have chosen to devote space on the focal plane to both guide and
Focus & Alignment (F&A) CCDs. F&A CCDs are mounted on the same plate as
the science focal plane array, but displaced vertically from that plane
by a distance sufficient to produce the classic "donut" image of an
out-of-focus star. With a known displacement, the measured size of the
donut allows a precise determination of the best focus position.
Analysis of changes in the shape of the donuts allow us to measure
misalignments between the DECam corrector and the primary mirror and
also provide an indication of low order aberrations derived from errors
in the primary mirror's figure and can be used to close the active
optics loop.

We will read out the F&A CCDs with each science exposure.

We expect \~25 well measured stars (S/N = 100) per science CCD in a 100
sec. science exposure. If the focus chips were 10 times smaller than the
science CCDs (i.e., focus chip area = 17 sq-arcmin), 2 stars per CCD
would be expected and we believe this to be sufficient. We assume the
guide exposures will be 1 second. One can measure sufficiently accurate
centroids to guide from a star of S/N=20. If we take the sky noise to be
16e- and the read noise over the star image to be 20e-, then we need a
star producing 1000 e- in 1 second. This is magnitude 19 for the *r* or
*i* bandpasses, which have surface densities of Σ = 0.2 per sq-arcmin
near the galactic pole. Assuming the stars are randomly distributed on
the sky, then the probability of finding no useable guide stars within a
guide camera of solid angle Ω is *e^-0.2×Ω^*. There are about 2000
fields in our survey. Guide CCDs covering 50 sq-arcmin yield a 10%
chance of the survey encountering a field with no guide stars and this
is an appropriate minimum areal coverage. In DECam, one 2kx2k CCD with
15 μm pixels covers 85 sq-arcmin.

### Guide CCDs

Guide CCDs will be sited within the dewar to the sides of the science
focal plane array.

The guide CCDs will observe the sky through the science survey filters.

32. Minimum guide CCD area: 20 sq-arcmin, conservative area: 50
    sq-arcmin

33. 

34. Guide CCDs should have 10e- read noise at 250 kpixel/s, 20e- at 500
    kpixel/s, readout in 1s in 4x4 binning

35. Guide CCDs require sub-array readout capability.

36. The guide CCD system must produce seeing measures suitable for use
    in the donut analysis as well as error signals suitable for guiding.

37. We require 4 2kx2k guide CCDs, spread on either side of the focal
    plane.

### Focus & alignment CCDs

Considerations of seeing, overlap and experience from other
observatories yields an optimum size for donut images of around 30
pixels, which results in:

38. F&A CCDs shall be displaced by 1.24mm perpendicularly to the FPA.

Donut size is affected by seeing, but this can be factored in by
concurrent seeing measurements from the guide CCDs (T-36).

Sufficient stars must be available for the donut analysis. Our
calculations and simulations of the star density we will encounter in
the survey lead us to:

39. F&A CCD minimum area: 20 sq-arcminutes, conservative area: 50
    sq-arcminutes

40. 

41. F&A CCD particulars same as science devices, read at same rate
    during survey

42. We require 4 2kx4k F&A CCDs, spread on either side of the focal
    plane.

All the required information can be derived from donut images taken on
one side of focus, but a useful opportunity for consistency checking is
provided by use of images from both sides of focus.

43. The four F&A CCDs should be arranged in pairs on opposing sides of
    the dewar on opposing sides of focus.

The F&A CCDs will be read out on the same cadence as the science FPA.
Updates to the corrector position via the hexapod and to the primary
mirror figure via the active optics system will be made at a cadence
designed for maximum benefit which will be determined by experiment at
the telescope.

The F&A CCDs will acquire stellar images through the survey filters.

## System Requirements for the DES optical corrector

Here we summarize additional technical specifications for the DES
corrector. We recognize that the camera will be used for a variety of
purposes by the general astronomical community and the corrector has
been designed with this in mind. Any changes in requirements beyond
those required for DES that have a significant cost impact will require
that those costs be borne by an organization other than the DES
Collaboration.

See section 3.4 for discussion of point spread function.

### Passbands:

44. The design shall be optimized for performance within individual
    passbands of 400-540, 560-680, 690-820, and 820-1000 nm.

45. It is allowed that the focus and/or scale may change among
    passbands.

46. Additionally, the corrector shall permit imaging in the UV (350 nm),
    although there are no formal specifications for throughput or image
    size.

The overall passband matches the optimal response of the LBNL CCDs. UV
capability is not required by the DES itself; however, the BIRP
committee explicitly expressed interest in knowing the capability in its
report.

### Ghosting

Two classes of ghosting need to be considered: re-imaging of the exit
pupil and ghost halos around bright stars. The former is endemic to
prime focus wide field correctors and causes problems with flatfielding
and photometric calibration. The requirements are:

47. Ghosting: The intensity of ghost image of exit pupil, expressed as
    the fraction of incident sky brightness reflect back as a ghost
    image, shall have a gradient over the length of one CCD (61 mm = 0.3
    degrees) of no more than 3%.

This is needed to achieve top-level photometric accuracy of 2%
(systematic error only). This requirement depends on the detailed
photometric calibration plan. As a baseline, we assume that flatfielding
is done using the sky background from median-averaged images, and that a
single zero-point (derived from standard stars) is used to convert ADU
to AB magnitudes. Ghosting introduced non-uniformities into the sky
background, creating artificial gradients that mimic sensitivity
variations. A gradient across a CCD introduces a) relative photometric
errors across the CCD; b) a photometric offset error relative to the
average of all stars on the device (depending on where the standard star
is positioned). The requirement is intended to allot 1/4 of the total
systematic variance to flatfielding.

If gradients are larger, one can correct for them using "star flats\";
however, these are tedious to obtain, and there will be some residual
errors from the process as well.

Note that steps in the ghost image are worse than linear gradients; star
flats then become necessary.

48. Ghost images of an in-focus star of 6 mag shall create a ghost image
    with a surface brightness no greater than 25 mag/sq-arcsec.

This derives from the need to achieve reliable detection and measurement
of faint objects (10*σ* detection at *r* = 23.4). Ghost images are
irregularly distributed across the focal plane and introduce local jumps
in sky brightness, causing error in measuring the local sky for
photometry of faint galaxies. (To second order, they affect the
construction of the median flat; we do not consider this here). We
expect one V = 6 star per DES field at latitude b = -50. The natural sky
brightness is about 21.5. The ghost image would be about 4% of this
brightness. This intensity is also comparable to what we expect in the
exit pupil ghost. An *r* = 23.4 galaxy of radius 100 has a surface
brightness of 24.6 mag/arcsec square. The local sky background should be
measured accurately for all except galaxies at the boundary of the ghost
image. In the worst case, a locally determined sky will be in error by
50%, causing a photometric error of 20%. This error is statistical in
form, since, when averaged over many DES fields, it is random in nature.

### Focal plane flatness

49. The optical design shall incorporate a flat focal plane. The peak to
    peak variation in matching CCDs to the focal plane shall be 30μm.
    The positioning of CCDs relative to a flat plane includes any CCD
    nonflatness.

It is easier to design and construct a mosaic array with a flat vs.
curved focal plane. Our studies show that no significant gains are to be
had by allowing the focal plane to be curved.

### Operating temperature range

50. The operating temperature for the corrector shall be -5 C to +27 C

These numbers are taken from CTIO website and reflect the annual
temperature variations experienced in the dome.

### Physical Size

51. The maximum diameter of any optical element shall be less than 1300
    mm.

52. The leading optical element shall be no more than 2300 mm in front
    of the focal plane.

These are the mechanical constraints placed by the available space in
the current design for the prime focus assembly. They do NOT include any
consideration of availability or associated risk for the size or cost of
any particular optical component.

### Glass Type

53. Glass type(s) shall be chosen according to availability, its ability
    to meet the design requirements, durability, low coefficient of
    thermal expansion, low radioactivity.

Fused silica seems to meet most if not all of these requirements.

54. BK7 is thought to be NOT acceptable for the dewar window due to
    radioactivity

### Throughput

55. Transmission shall be uniform in wavelength (≤10% peak-peak
    variation) over the full wavelength range specified in requirement
    (S-5; note also section 1.1.1 and requirement T-46.)

We do not want to compromise performance due to the choice of a glass
that would reduce transmission at either extreme of wavelength. (It is
thought that optical quality fused silica would meet this requirement.).
The same is to be said about the choice of anti-reflection coating,
though the issue here is uniformity of the coating over the glass.

The number 10% is a soft target, but if it exceeds this value then we
must evaluate it in the context of the spatially varying system response
requirements.

### Filters

DECam must accommodate community use and this will require the facility
to use filters other than those specified by the DES.

56. A filter of at least 15 mm thickness shall be accommodated in the
    design. If possible, the filter shall be placed at a location that
    minimizes its diameter. The filter changer mechanism shall
    accommodate at least 6 filters.

The filter must be thick enough to be mechanically sturdy. Sag due to
gravity is thought to be immaterial on the optical design for this
thickness.

Displacement of filters between exposures can result in photometric
errors from the movement of transmission inhomogeneities with respect to
the FPA.

57. The filter mechanism must place each filter within the corrector to
    a precision of 1mm or better.

### Filter transmission

The signal to noise calculation of section 3.6.1 also considered filter
transmission, the requirements are given in Table 3.7.

58. Filter transmission requirements: 85% in \[*g, r, i, z*\]

Table 3.7 Filter Transmission Requirements

  --------- ------------------ -------------------- ---------------------
  filter    CWL (nm)           FWHM(nm)             Transmission

  *g*       475                150                  85%

  *r*       635                150                  85%

  *i*       775                150                  85%

  *z*       925                150                  85%
  --------- ------------------ -------------------- ---------------------

### Dewar window

59. 

60. 

61. Dewar window and wall thicknesses will comply with all appropriate
    safety standards to ensure the integrity of the instrument. The
    current design calls for a window thickness of 65mm to minimize the
    risk of catastrophic failure through inadvertent impacts.

62. The dewar window may have power.

63. The curvature of the window induced by the air-vacuum pressure
    differential must be computed and its impact on the optical design
    assessed.

It is thought that the sag could be of order 60 microns, which would
require such an assessment.

64. The dewar window glass type may be specified as part of the design.
    The material shall be chosen for low radioactivity. Potassium
    content shall be no more than 1%, and trace thorium that be no more
    than 10 ppm by weight. This excludes BK7, which has 10 percent
    potassium (requirement T-54).

Radioactivity generates spurious charge events in the CCD and must be
minimized. The radioactive isotopes one might find in glass include K-40
and Th-232. Inside the vacuum, beta rays from radioactive decay on the
surface of the dewar window can strike the CCD; such particles will be
absorbed and produce complex tracks (unlike cosmic ray muons, which pass
through the CCD on straight-line trajectories). The upper limits
represent the amount of material of each type that will produce an event
rate about equal to that of cosmic rays at sea level.

### Shutter

Placing the shutter closer to the focal plane means a smaller beam size.
Space for the shutter is constrained in the prime focus cage. The
location is not critical but should be as close to the focal plane as
feasible.

Shutter performance requirements are based on the precision required to
make calibrations via standard star observations:

65. The shutter should provide 1 percent uniformity at a 1 second
    exposure time.  (i.e. the actual exposure time anywhere on the CCD
    should not be more than 1% different from anywhere else at this
    exposure time).

66. Shutter exposure time repeatability should be better than 0.005
    seconds (5 milliseconds).

67. Shutter exposure time accuracy should be such that the offset from
    the nominal exposure should be less than 0.05 seconds (ie 50
    milliseconds)

68. Absolute timing of an exposure should be measured to a precision of
    10ms and recorded in the resulting image FITS header.

69. The shutter should allow exposures from 1 second upwards.

We will not test performance for shorter exposure times, but capability
to take shorter exposures (e.g. down to 0.2 seconds) would be useful as
goal, not a specification.

70. Shutter performance should be maintained from -5°C to +20°C.

71. Shutter performance should be maintained after high duty-cycle
    tests.

72. Approximately 125,000 shutter movements will take place per year,
    Shutter MTBF and MTTR should be based on a 10-year lifetime.

Monitoring and diagnosis of performance anomalies (uniformity,
repeatability, accuracy above) would be aided by suitably telemetry from
sensors associated with the shutter mechanism.

### Cost

73. Within the constraints of the requirements on performance, the
    optical design shall minimize the total cost of construction.

Costs are driven by the glass type, spherical vs. aspherical surfaces,
details of lens shape (e.g. thickness), the total number of elements
(both with respect to procurement and to mounting complexity), the
delivery schedule, and the availability of qualified vendors.

## Instrument Cage Motion

This section discusses requirements to apply to any mechanism for
adjusting the position and orientation of the instrument cage. X and Y
refer to motions of the corrector and camera in directions perpendicular
to their optical axis. "Tilt" refers to rotations of the corrector and
camera about the X or Y axes. "Focus" refers to motions of the corrector
and camera in the direction parallel to their optical axis. "Rotation"
refers to rotations of the corrector and camera about their optical
axis. The hexapod will provide control over all these degrees of
freedom.

The total error budget for fabrication, assembly, etc is 0.1" FWHM (=
3.8 μm rms radius). It is assumed that this is contributed by 100
degrees of freedom Thus, for any single degree of freedom, we assume an
error budget of 0.01" FWHM or 0.38 μm rms radius, yielding the
tolerances given on focus, X or Y or X-tilt or Y-tilt.

A change in X of 45 microns introduces a mean whisker of 0.06 arcsec. A
change in focus of 7 microns introduces a mean whisker of 0.06 arcsec.
Note that whisker length increases as sqrt(offset).

Note that rms radius increases linearly with offset.

### Focus

74. The tolerance on setting the focus is +/- 7 microns peak-peak error.
    Thus, the focus resolution should be 1 microns.

75. Range of focus motion should be at least 4 mm to accommodate the
    temperature extremes.

A change of 30 C corresponds to 3.6 mm; (one must be sure to position
the corrector properly within the adjust range!) If the range is 8 mm,
one can be cavalier with the initial installation of the corrector -
just put it in the middle of the range.

The actual performance of the focus loop needs to be better than SDSS
achieves. The SDSS can detect a defocus at the focal plane of 12
microns. Scaling by telescope f/ratio, we should be able to detect a
defocus of 7 microns at the Blanco.

### Tilt and X&Y translation

76. We assume a worst-case de-center of the primary mirror with respect
    to the corrector is 1.5 mm.

77. Only translation or tilt of the corrector assembly is needed - not
    both. In practice, tilt is not feasible, if the pivot point is in
    the cage.

Correction of a translation of the primary mirror by a tilt of the
corrector yields a residual tilt in the focal plane of +/- 1 micron at
the edges, which is not measurable.

#### X&Y Translation

78. The range of lateral translation of the corrector should also be +/-
    1.5 mm in each of X and Y directions.

79. The tolerance in setting the corrector lateral translation is +/- 50
    microns peak-peak error. The step size should be 10 microns.

#### Tilt

80. The range of corrector tilt (if provided) shall be +/- 30 arcsec in
    each of the X and Y directions.

81. The tolerance on setting the tilt of the corrector is +/- 1 arcsec
    (2.8×10^-4^ degrees) peak-peak. The step size should be 0.2 arcsec.

82. If no tilt correction is provided, the tilt angle shall be
    maintained to an accuracy of 2 arcsec over the full range of
    translation.

### Orientation changes

Because of the top-end flip capability of the telescope, DECam will
experience all possible orientations. Even when DECam is not in use, it
will be necessary to run any and all instrument components for
diagnostic and test procedures.

83. All instrument components shall perform to specification at all
    possible orientations.

84. The maximum change in telescope zenith angle during one exposure
    shall be 5 degrees.

This angular change specification corresponds to a 20 minute exposure
time at the most unfavorable orientations. If needed, a more detailed
chart can be prepared giving angular change vs. hour angle and
declination.

##  Summary of Technical Specifications

The full list of technical specifications is as follows:

1.  Number of tilings - The minimum number of tilings of the survey area
    in any filter is 4.

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85. Integration Times -- \[*g, r, i, z*\] = \[400, 400, 1200, 2000\]
    sec.

86. Single exposure Time -- minimum 100 sec

87. Observing time contingency: 25%.

88. Pixel scale: 2.5 pixels per FWHM, met by 0.27"/pixel, 15μ

89. The delivered PSF from the telescope, instrument, and other factors
    exclusive of the site seeing shall be no greater than 0.55 arcsec.

90. No requirement is placed on the static PSF shape.

91. For each dynamic source of systematic error (e.g., focus
    variations), the peak amplitude induced in the PSF whisker length
    will be no larger than 0.06".

92. Charged particle events due to radioactive material in the
    instrument shall produce a background no worse than that arising
    from cosmic rays at sea level (0.01 events/cm^2^/sec).

93. CCD nonlinearity: 1% or better

94. CCD full well: 130,000e-

95. CCD residual image: ≤ 10e- from 3×10^6^ e- spread over 5 pixels

96. CCD readout rate: 250 kpix/s

97. CCD dark current: ≤ 0.007e-/s/pix

98. CTE: \> 0.99999

99. CCD blooming control: not required

100. CCD prescan pixels: Not required, but acceptable

101. CCD frame transfer: not required

102. Outputs/CCD: 2

103. CCD quantum efficiency -- \[*g, r, i, z*\] : \[60%, 75%, 70%, 65%\]

104. CCD QE stability: ≤ 0.3% variation in QE over 12-18 hours

105. CCD temperature stability - 1 degree K stability over 12-18 hours.

106. Spatial temperature variation across the focal plane should produce
     ≤ 5% variation in QE

107. Spatial temperature variation: ≤ 15 degrees K across the focal
     plane.

108. CCD read noise must be ≤ 15e-/pixel

109. Goal for CCD read noise at DES readout rates: 10 e-/pixel

110. The CCD charge diffusion should have an rms 1-D width *σ* no
     greater than 7.5μm.

111. CCD flatness: on \< 1 cm^2^ scales on a given CCD, the mean
     flatness variation should be ≤ 3μm.

112. CCD flatness: between adjacent ≤ 1cm^2^ areas on a given CCD, the
     difference in mean flatness should be ≤ 10μm.

113. Cosmetic defects should account for no more than 5% loss in CCD
     area from non-usable pixels.

114. Guide error signal update rate: 1 Hz.

115. Minimum guide CCD area: 20 sq-arcmin, conservative area: 50
     sq-arcmin

116. Guide CCDs should have 10e- read noise at 250 kpixel/s, 20e- at 500
     kpixel/s, readout in 1s in 4x4 binning

117. Guide CCDs require sub-array readout capability.

118. The guide CCD system must produce seeing measures suitable for use
     in the donut analysis as well as error signals suitable for
     guiding.

119. We require 4 2kx2k guide CCDs, spread on either side of the focal
     plane.

120. F&A CCDs shall be displaced by 1.24mm perpendicularly to the FPA.

121. F&A CCD minimum area: 20 sq-arcminutes, conservative area: 50
     sq-arcminutes

122. F&A CCD particulars same as science devices, read at same rate
     during survey

123. We require 4 2kx4k F&A CCDs, spread on either side of the focal
     plane.

124. The four F&A CCDs should be arranged in pairs on opposing sides of
     the dewar on opposing sides of focus.

125. The design shall be optimized for performance within individual
     passbands of 400-540, 560-680, 690-820, and 820-1000 nm.

126. It is allowed that the focus and/or scale may change among
     passbands.

127. Additionally, the corrector shall permit imaging in the UV (350
     nm), although there are no formal specifications for throughput or
     image size.

128. Ghosting: The intensity of ghost image of exit pupil, expressed as
     the fraction of incident sky brightness reflect back as a ghost
     image, shall have a gradient over the length of one CCD (61 mm =
     0.3 degrees) of no more than 3%.

129. Ghost images of an in-focus star of 6 mag shall create a ghost
     image with a surface brightness no greater than 25 mag/sq-arcsec.

130. The optical design shall incorporate a flat focal plane. The peak
     to peak variation in matching CCDs to the focal plane shall be
     30μm. The positioning of CCDs relative to a flat plane includes any
     CCD nonflatness.

131. The operating temperature for the corrector shall be -5 C to +27 C

132. The maximum diameter of any optical element shall be less than 1300
     mm.

133. The leading optical element shall be no more than 2300 mm in front
     of the focal plane.

134. Glass type(s) shall be chosen according to availability, its
     ability to meet the design requirements, durability, low
     coefficient of thermal expansion, low radioactivity.

135. BK7 is thought to be NOT acceptable for the dewar window due to
     radioactivity

136. Transmission shall be uniform in wavelength (≤10% peak-peak
     variation) over the full wavelength range specified in requirement
     (S-5; note also section 1.1.1 and requirement T-46.)

137. A filter of at least 15 mm thickness shall be accommodated in the
     design. If possible, the filter shall be placed at a location that
     minimizes its diameter. The filter changer mechanism shall
     accommodate at least 6 filters.

138. The filter mechanism must place each filter within the corrector to
     a precision of 1mm or better.

139. Filter transmission requirements: 85% in \[*g, r, i, z*\]

140. The dewar window may have power.

141. The curvature of the window induced by the air-vacuum pressure
     differential must be computed and its impact on the optical design
     assessed.

142. The dewar window glass type may be specified as part of the design.
     The material shall be chosen for low radioactivity. Potassium
     content shall be no more than 1%, and trace thorium that be no more
     than 10 ppm by weight. This excludes BK7, which has 10 percent
     potassium (requirement T-54).

143. The shutter should provide 1 percent uniformity at a 1 second
     exposure time.  (i.e. the actual exposure time anywhere on the CCD
     should not be more than 1% different from anywhere else at this
     exposure time).

144. Shutter exposure time repeatability should be better than 0.005
     seconds (5 milliseconds).

145. Shutter exposure time accuracy should be such that the offset from
     the nominal exposure should be less than 0.05 seconds (ie 50
     milliseconds)

146. Absolute timing of an exposure should be measured to a precision of
     10ms and recorded in the resulting image FITS header.

147. The shutter should allow exposures from 1 second upwards.

148. Shutter performance should be maintained from -5°C to +20°C.

149. Shutter performance should be maintained after high duty-cycle
     tests.

150. Approximately 125,000 shutter movements will take place per year,
     Shutter MTBF and MTTR should be based on a 10-year lifetime.

151. Within the constraints of the requirements on performance, the
     optical design shall minimize the total cost of construction.

152. The tolerance on setting the focus is +/- 2.5 microns peak-peak
     error. Thus, the focus resolution should be 1 microns.

153. Range of focus motion should be at least 4 mm to accommodate the
     temperature extremes.

154. We assume a worst-case de-center of the primary mirror with respect
     to the corrector is 1.5 mm.

155. Only translation or tilt of the corrector assembly is needed - not
     both. In practice, tilt is not feasible, if the pivot point is in
     the cage.

156. The range of lateral translation of the corrector should also be
     +/- 1.5 mm in each of X and Y directions.

157. The tolerance in setting the corrector lateral translation is +/-
     50 microns peak-peak error. The step size should be 10 microns.

158. The range of corrector tilt (if provided) shall be +/- 30 arcsec in
     each of the X and Y directions.

159. The tolerance on setting the tilt of the corrector is +/- 1 arcsec
     (2.8×10^-4^ degrees) peak-peak. The step size should be 0.2 arcsec.

160. If no tilt correction is provided, the tilt angle shall be
     maintained to an accuracy of 2 arcsec over the full range of
     translation.

161. All instrument components shall perform to specification at all
     possible orientations.

162. The maximum change in telescope zenith angle during one exposure
     shall be 5 degrees.

[^1]: At the time of commissioning there will be additional questions:
    iv) what are the tests, v) who is doing them, vi) when are they to
    be done, and vii) with what resources. These we leave for later.