Other resources you may find useful are listed below; they are all available via the Kitt Peak documentation page on the web at: http://www.noao.edu/kpno/docs.html.
|SITe||Mosaic||2048 x 4096||15||70,000||6.0||5.0||?||3|
|4m+Mosaic||3.1||36 x 36||0.26||left||down|
|WIYN+Mini-Mo||6.5||10 x 10||0.141||left||down|
|2.1m+T2KA||7.5||10.2 x 10.2*||0.305||left||down|
|0.9m+Mosaic||7.5||59 x 59||0.43||left||top|
|*Unvigenetted FOV at the 2.1-m is 2048 x 1850 pixels, or 10.2 x 9.4 arcminutes, with the vignetted region being to higher column numbers (south).|
|The relative DQEs of the CCDs used for direct imaging are shown above. The SITe DQE curve is typical of the Mosaic and Mini-Mosaic chips.|
The standard UBVRI broad-band set is kept with each camera. In addition we have numerous other filters available in either 4in x 4in or 2in x 2in size, including a Gunn set, a Stromgren set, and a Washington system set, as well as numerous interference filters for narrow-band work. A complete listing of filters, and copies of their transmission plots is available via the Kitt Peak home page.
In selecting narrow-band filters for your observing run, keep in mind that there is a significant shift of the central bandpass to the blue (by as much as 20Å) when using interference filters in fast f-ratio beams. The transmission of a narrow-band filter in these fast beams can be simulated; if you need information about a particular filter please contact us via "email@example.com".
Users of Mosaic should refer to the Mosaic
manual; similarly, users
of Mini-Mo should refer to the Mini-Mosaic
manual for focusing
procedures. Typical focus values and offsets for our commonly used
filters are given in the table below, along with recommended step
sizes to use when focusing. The focus offsets for other 4x4-inch
filters, is given later in this section. By interpolating you will
likely be able to determine the best focus to half of one of these
steps or better; this has been made particularly easy for the user
with the IRAF task kpnofocus (see the ICE manual).
There is a fairly strong focus dependence on temperature at all telescopes; in addition, the 2.1-m shows a pronounced effect with zenith distance. Once you have a good focus, the coefficients below will guide you in refocusing as the temperature and airmass changes.
Thus the focus for the BVRI filters at the 2.1-m can be predicted as:
|F ~ 20500 - 75 x (T-10) -130 (X-1.0)|
where the temperature T is measured in oC and X is
Guiding and Acquisition
All of the telescopes used for direct imaging have automatic guiders
that should be used for exposures longer than a few seconds. The
following is a brief description of the guider characteristics. At
the 4m and WIYN the telescope operators handle setting up for
guiding; complete details of how to use the guiders at the 2.1-m
and 0.9-m can be found in the
2.1m Observer Handbook
WIYN 0.9m Observing Manual.
Note that in the case of all of the shutters that the entire field is illuminated for the identical time: these shutters are of "guillotine" type, with the blades moving the same way on opening and closing for one exposure, and then the opposite way on the following exposure.
|WIYN Mini-Mo||0.000 ±0.002|
If you are attempting to do 1% stellar photometry of stars in a cluster, you are probably interested in covering as large a magnitude range as possible, and furthermore, your noise is going to be primarily photon-noise, not read-noise. Go for the largest value of e- per ADU as you can without exceeding the linearity of the particular chip. Generally, this will be the default gain, which is also the gain that will give you the least amount of horizontal bleeding from very saturated stars and the most uniform noise characteristics.
If you are doing surface brightness studies of objects through narrow-band filters, and the read-noise is significant but the dynamic range of your objects is limited, you may wish to stay with the largest gain number (smallest number of e- per ADU). Similarly in some very low-signal spectroscopic applications you are limited by the read-noise.
Note: T2KA shows some very low-level "streaking" to the right
of the most heavily exposed stars. The electronics has been adjusted
to minimize this problem for the default gain setting. If you are
concerned about the effects of very saturated stars, you would do
well to stay with the default gain setting.
Dealing with Cosmic Rays
If you were to take a 15-minute "dark" exposure you would be struck
by the large number of 1-2 pixel radiation events present. Fortunately,
most of these are of modest amplitude and would be lost in the
sky-noise for broad-band work. However, some would not. These
"radiation events" are often dubbed "cosmic rays", although in
fact many of them are secondaries that originate in the CCD substrate
In order to filter out these radiation events most observers will
divide their long exposures (>10 min) into 2 or 3 pieces. With
the low read-noise present in our CCDs, combining three 10 minute
exposures usually has the same signal-to-noise as a single 30 minute
exposure. The only loss is the extra read time.
By "calibration data" we refer both to what is needed to remove
the "instrumental signature" of the CCD (biases and flat-fields,
say) as well as what is needed to provide photometric and/or flux
What Calibration Data You Need and Why
The goal in removing the instrumental signature is to transform the
data so that the output signal is linearly proportional to the
amount of light entering the telescope, at least as linear as the
CCD allows. To do this one needs to first remove the additive terms
(such as bias structure and dark current) and then remove both the
pixel-to-pixel gain variations and the lower frequency response of
the telescope/detector to a uniform (flat) level.
Over-scan and bias: With every exposure there is a pedestal level which is added to the output signal: "the bias", typically several hundred ADUs. As the temperature of the electronics changes during the night this bias level will also change by a few ADUs. All of our chips are read with an extra 32 columns of "overscan" which provides a measurement of this bias level. In data reductions one can use this strip to determine either a scalar correction, or fit a smooth function to the level as a function of line number.
Some chips exhibit a spatial bias structure as well, and it is necessary, in those cases to use zero-second bias frames to correct for this (constant) two-dimensional structure. By using 10 or so frames to construct an average one obtains a calibration frame which will not increase the noise much on one's data frame.
One thing that you may notice on your bias-frames is "banding"---these are usually low-amplitude (1-3 ADU) bands that extend across all columns (including the over-scan) and are perhaps 20-100 lines in width. These bands are due to slight electronic noise, and unfortunately will vary with every exposure. One has two recourses: (a) either one can ignore these (for most direct imaging applications these bands will be completely lost in the sky-noise), or (b) one can remove these by using a high-order fit in the overscan (cubic-spline of order 100-200, say). Although these may appear to be quite severe in the bias frames, one should measure their amplitude before panicking. If the amplitudes are appreciable (>10 ADU, say) then these are probably curable and a call for help is warranted.
Dark-current: All CCDs suffer from thermal "noise" as electrons jitter around and are occasionally liberated; these non-photon events are then trapped in the potential well. To reduce the size of this dark-current our CCDs are cooled to approximately -100 oC with the result that dark-current on our chips is barely detectable (3-4 e-/hr/pixel). The dark-current is usually quite uniform and hence has no effect on photometry if the sky value is being determined from the program frames themselves. One might still wish to take several dark-frames to substantiate that the effect is small. In practice, light-leaks can exceed the level of the dark current level, so care should be done to do this with the dome darkened. The exposure time on your dark-frames should equal or exceed your longest exposure time, and if there is any possibility that you may actually need to use these dark-frames to correct your data, you should take a minimum of 3 of these to allow cosmic-rays and radiation-events to be filtered out in determining the average.
Flat-fields: Each pixel in the CCD responds to light a little differently. Modern CCDs are surprisingly "flat" in their response (i.e., the gain of each pixel is nearly the same), but as Mackay (1986 ARAA, 24, 255) put it: "The only uniform CCD is a dead CCD." In practice there are slight color-dependent gain differences between each pixel that must be removed for good photometry.
To remove this pixel-to-pixel gain variation one needs a series of well-exposed "flats", obtained through each filter. The total exposure in one's flats should be such that one never degrades the signal-to-noise in your program frames; in practice, accumulating 4 or 5 flats with 20,000 e-/pixel will amply suffice to remove the pixel-to-pixel variation. Better matching of response to very low illumination levels could require a more extensive series of exposures at an even lower fraction of full well.
However, even if the telescope were being illuminated completely uniformly (by, say, the night sky) the CCD is unlikely to be illuminated uniformly. Instead, vignetting due to guider mirrors, non-uniformities in the filters (and dust on the filters) require that large-scale flat-fielding is necessary.
How to best achieve this appears to be telescope-dependent on Kitt Peak, and we give our recommendations in the section Recommendations. Possibilities include:
Observations of standard stars permit removing the effects of atmospheric extinction and transforming your instrumental magnitudes to "standard magnitudes" and/or flux. Since the scientific goals of some programs require only good relative photometry, while others ask for (and should be able to achieve) 1% all-sky photometry, we can offer you only general guidelines in obtaining adequate photometric calibration of your data. The definitive and best-calibrated set of UBVRI photometric standards are those of Landolt (1973, AJ 78, 959; 1983, AJ 88, 439; 1992 AJ 104, 340); copies of these papers, including finding charts, are kept at each telescope in a single binder.
Most observers' choice for flux calibration of narrow-band images are the
"Kitt Peak Spectrophotometric Standard" stars found in Massey et al (1988,
ApJ 328, 315) and Massey et al (1989 ApJ 358,
344) or the KPNO-produced IRS and IIDS Standard Star Manuals. There are
caches of the coordinates of the Landolt and Spectrophotometric Standards
on the computers at the 4-m, 2.1-m.
Obtain 20 of these each afternoon and compare the combined,
overscan-corrected frame to that from the previous day. Daily bias
frames provide a good check that the instrument is performing
Dark frames: Obtain a 15-minute dark-frame with the dome as dark as possible. Process it and examine for excess counts. If your program requires dark frames, obtain a minimum of 4 over the course of your run, each with an exposure time equal to your longest exposure time.
Flat-fields: The procedures for obtaining good flat-field exposures are given below, followed by our telescope-specific recommendations of what you may need.
Dome-flats: Obtain 5 or more of these (each 20,000 e- per pixel) through each filter at the white spot. For BVRI use the lamps with the blue, "color-balance" filters; for U or for narrow-band interference filters, use the lamps without the color-balance filters. The table below will give an approximate guide to exposure times, but by all means check these by first doing a "test" exposure with each filter to substantiate that the exposures are neither saturated nor underexposed. (Note: "doobs" is available for running through such a sequence automatically if the lamps do not have to be adjusted; see the Mosaic Manual, Mini-Mosaic Manual, or the ICE manual.)
Note that dome flats run during the daytime (particularly at the 0.9m) may be seriously compromised due to daylight leaking into the dome and straying into the telescope. We recommend that before you run your dome flats you try an exposure with the appropriate exposure time with the flat-field lamps off This will tell you whether or not scattered light will contribute significantly. You will have to try this at various wavelengths, as the scattered light problem can be especially bad at R and I
|Filter||4-m Mosaic||WIYN+Mini-Mo||2 Meter+T2KA||0.9-m+Mosaic|
Twilight-flats: At the 4 meter Mosaic, as soon after sunset as possible obtain twilight flats with the telescope tracking and the dome clear. For broad-band work, begin with your bluest filter and wait until a one-second exposure no longer saturates the chip. Obtain 3-4 exposures through each filter, stepping the telescope by 30 arcsec between each exposure, and changing the exposure time as needed to avoid underexposing or saturation. It is possible but difficult to obtain a complete set (UBVRI) during a single twilight, so this may take several evenings and mornings to obtain a complete set. Twilight flats obtained during partially cloudy conditions do not appear to work. Twilight flats do not appear to work at either WIYN or the 2.1-m, both open-tube telescopes, probably due to the large amount of scattered light present in the dome during bright twilight.
Dark-sky flats: If your program frames are relatively sparse and contain sky levels at least 100e- above bias, you can probably use these frames themselves as the ultimate correction for large-scale illumination errors. You will have to combine them with scaling by the mode and experiment with the various rejection algorithms to remove all your interesting stuff (stars/galaxies) and leave you only sky. Alternatively, you can obtain 4-5 exposures (the more the better) of relatively "blank sky" fields, offsetting the telescope by 30 arcsec or more between each exposure. When done you will be able to "clip" out the stars and smooth the result. Still, you will want at least 100 e- per pixel above sky if you want to use these to correct to a small fraction of the night sky. Although stars appear to be just about everywhere, the following fields are not as full of them as others:
Fringe Frames: However obtained, your dark-sky flats will also suffice for removing the average fringing in R and I. However, as the relative intensities of the night sky lines that cause fringing in the I band are quite variable during the night, don't expect to do a perfect job. Remember that fringing from night-sky lines can be treated as additive and that an appropriate scale-factor must be determined. See the "User's Guide to Reducing CCD Data With IRAF" guide.
What Works/What Doesn't
These numbers are not from stone tablets, but you should see similar count rates to within 25% or so.
How long do you have to integrate to obtain a given signal-to-noise ratio (SNR) at some magnitude? The SNR will simply be
N=count rate in electrons per second per image.
R=Read-noise in electrons; this can be taken to be 4 e- for all practical purposes (see CCD Characteristics.
p=Number of pixels in stellar image. If the seeing profile FWHM of r (in pixels), then it is not unreasonable to use an area of
What do we use for the sky brightness? The table below is taken from Alistair Walker's article on the sky brightness at Tololo in NOAO Newsletter No. 10. Photometry of the Kitt Peak night sky is consistent with these numbers.
|lunar age (days)||U||B||V||R||I|
We can easily turn this into a count in electrons per second per pixel (e.g., "S" in the above equation by simply using the telescope/chip sensitivities given above and the scale (arcsec per pixel) given in CCD Characteristics. The table below gives this for new moon; other lunar phases can thus be readily computed using the sky brightness relative to new moon.
Thus the (very rough!) expected count rate from the sky in V at WIYN with Mini-Mo at full moon will be 10(21.8-20.0)/2.5 × 2.0 = 10.5 e-/sec/pixel.
We can attempt to duplicate the result given in our "Limiting Magnitude" table: What is the integration time required at WIYN and S2KB obtain a SNR of 50 at V for a star of 23.6 magnitude during with a Quarter Moon (lunar age=7 days) in the sky? We expect the answer to be about an hour!
The count rate N for a V=23.5 mag star will be 10(20.0-23.6)/2.5 × 270 = 9.8 e-/sec/image.
The sky counts S will be 10(21.8-21.4)/2.5 × 1.0=1.5 e-/sec/pixel.
The number of pixels will of course depend upon the seeing; if it is something like 1.0 arcsec then the number of pixels will be 1.4 × (1.0/0.14)2 = 71 pixels.
We then have:
A=N2 = 9.82 = 96
B= -(SNR)2× (N+pS) = -502 × (9.6+71× 1.5) = -290250
C = -(SNR)2 × pR2 = -502 × 71 × 62 = -6390000
and the time will be:
t = (-B + sqrt(B2-4AC))/2A = 3050,
close enough for government work.
Color Transformation and Extinctions
The following are "typical" values from a run on the 2.1-m with
T1KA. The color terms may be considered typical of our thinned,
AR-coated, backside-treated CCDs (e.g., T2KA). If you have recently
done a color transformation, let us know!
In the above, lower-case are the instrumental magnitudes,
zero-point constants, and X is the airmass. Don't be surprised
if your extinction terms vary by 0.1 mag from the values listed,
although such variations are typically grey (i.e., all coefficients
vary by a single additive value from those above).
Standard UBVRI Filters
Our "standard" UBVRI filters are collectively known as the
"Harris" set. Historically the B and V filters are glass
sandwiches designed by Hugh Harris (USNO) in order to yield small
color-terms. The R filter is a glass sandwich designed by Alistair
Walker (1986, IAU Symp. 188, p. 33) to match the Kron-Cousins
system. The I filter is an interference filter designed by Jeremy
Mould. Our U filter is a liquid CuSO4 plus UG1 glass sandwich,
the origin of which has been lost in the mists of time. (Schott
quit making UG2, so UG1 is used in its place in the newer filters;
the difference is very slight.) Specifically our filters are
constructed as follows:
Note that our UBVRI filters are not all quite identical; tracings of individual filters reveal differences as great as 70 Å in the central wavelengths.