Changes from EDR to DR1

The SDSS data, both imaging and spectroscopy, are processed through an extensive series of software pipelines, which have been developed over a number of years. The versions of the software used for the Early Data Release (EDR) met most (although not quite all) of our scientific requirements. This software and its performance is described in detail in Stoughton et al (2002) (hereafter, EDR paper), where a number of known problems and caveats are laid out. Since that time, there have been substantial revisions to all the pipelines which have been used to process the DR1 data; this document describes the most important of these changes and the effect that these have on the science outputs of the survey. Remaining known problems are described as well. A summary description of DR1 is given in the DR1 paper (Abazajian et al 2003).

Photometric Calibration

As described in the EDR paper (sections 3.2.1 and 4.5), our photometric calibration system is complicated by the fact that we calibrate with a separate telescope (the 20-inch 'Photometric Telescope', or PT) from the one with which the imaging data are taken (i.e., the 2.5m). Because of unfortunate properties of the filters used in these observations, the effective band-passes on the PT and the 2.5m differ systematically. Our previous formalism for carrying out the photometric calibration took this into account in an approximate, and not completely self-consistent way. We have now reformulated the photometric equations, so that we work directly in the natural system of the 2.5m, making the relation between detected counts and magnitude a simple one. The exact details are described in this definition of the SDSS photometric system (.ps, 80 kb). (See also this webpage for a summary of the photometric equations now in use.)

One consequence of making the photometry self-consistent is that we no longer have to indicate our photometry as preliminary with the asterisk notation. Thus photometry as measured by the 2.5m, and as released as part of DR1, should be referred to as u g r i z (and not u* g* r* i* z*). Note that the mean colors of stars on the old and new systems (i.e., u*g*r*i*z* vs. ugriz) have been forced to be the same.

Conversion to AB magnitudes is difficult

As Section 4.5 of the EDR paper discusses, putting these magnitudes on a consistent AB system (i.e. converting the magnitude zeropoint into physical flux units) is tricky. We expect that our u g r i z photometry is not exactly on the AB system, but we estimate from several lines of argument that the g-r, r-i, and i-z colors are nearly AB to within about 3%. We estimate that the u-g color is too red by about 5% (again with about 3% of uncertainty), in the sense that (u-g)(AB) = (u-g)(SDSS) - 0.05. We are continuing to work on measuring and cross-checking the exact photometric AB zeropoints. Note that our relative photometry is quite a bit better than these numbers would imply; repeat observations show that our calibrations are better than 2%.

The Photometric Pipeline

This is the code that flat-fields and bias subtracts the images, corrects for cosmic rays and bad columns, and finds and measures the properties of objects. As described in Section 4.5.4 of the EDR paper, the flat fields are affected by scattered light, especially in the u-band, thus giving systematically incorrect photometry. We have now characterized the flat fields much more accurately, and reduced this systematic effect to be no larger than 2% in g, r, i and z, and 3% in u. RMS statistics are substantially better; uncertainties in the flat fields contribute less than 1% rms to the photometric error in u, and less than 1/2% in all other bands.

Improved PSF fitting means improved photometric calibration

Much of the photometric calibration depends on a detailed understanding of the Point Spread Function (PSF); both its width as a function of position and time, and its outer extent in order to carry out an aperture correction (EDR paper Section 4.4.5.2). The old version of the pipelines had trouble following these variations when the PSF was changing rapidly; this code has been made substantially more robust now, and should remove residual effects in the photometry as a function of seeing. It is still true that there are small regions of the data in which the seeing variations are so rapid (FWHM changing by more than 0.2 arcsec in a minute) that the current code will have systematic errors, but these errors are estimated and propagated through to the quoted errors in the PSF photometry.

Non-linearity correction applied

There is a small but measureable non-linearity in the response of the photometric CCD's, measuring several percent at saturation. This is now corrected for explicitly. This is important, as the photometric calibration is tied down at the bright end, where this non-linearity can matter. See the non-linearity documentation for details.

Improved cosmic-ray rejection

Cosmic rays are recognized as such by their sharp gradients relative to the PSF. This is effective in removing the vast majority of them but, as described in Fan et al (2001), is not perfect, especially when looking for rare objects detected in only one band. That paper describes an enhanced routine to look for signs that a given detection might be a cosmic ray masking as an object; this routine is now run as part of PHOTO. A flag, OBJECT2_MAYBE_CR, is set for objects which are suspected to be cosmic rays.

Deblender improvements

The deblender (EDR Paper 4.4.3) separates overlapping images of galaxies and stars. It had the unfortunate behavior of often shredding large galaxies with substructure into several pieces, especially brighter than r ~15 mag. This behavior has been suppressed, and the vast majority of bright galaxies are now treated properly by the deblender. The deblender works by noticing peaks in the image and assigning to each a template shape from which to carry out the deblending. If two templates were substantially the same (i.e., if they were close to degenerate), the deblender tended to give unphysical results. This was fixed by removing from the list of templates those that were close to degenerate. The photo flag OBJECT2_DEBLEND_DEGENERATE is set to indicate parents for which this template pruning has happened. A related useful flag is OBJECT2_BRIGHTEST_GALAXY_CHILD, which is set for that child in a complicated deblend that is probably associated with the bright galaxy in question (as the name implies). See the deblending flags information.

Adaptive moments added

In addition to the object shape measures described in Section 4.4.6 of the EDR paper, the photometric pipeline now calculates so-called adaptive moments which weight the photons in a close-to-optimal way for weak lensing measurements of faint objects. These are described in detail in the Adaptive Moments algorithm page. Briefly, the quantities calculated include:

M_e1 = (Mxx - Myy)/(Mxx + Myy)
M_e2 = 2 Mxy/(Mxx + Myy)
M_rr_cc = Mxx + Myy

Four items of the covariance matrix are stored. Each term as listed is the square root of the corresponding entry in the covariance matrix; for negative covariance terms, the quantity sign(x)*abs(sqrt(x)) is stored. The terms are:

M_e1e1Err: The square root of the e1 variance
M_e2e2Err: The square root of the e2 variance
M_e1e2Err: The square root of the e1-e2 covariance (with the sign convention above)
M_rr_ccErr: The square root of the rr-rr covariance

M_e1_psf
M_e2_psf
M_rr_cc_psf
M_cr4_psf

Various flags are set indicating possible problems in determining these moments. A complete description of the PHOTO flags is given in the description of photo flags (see the description of moment-related flags).

Improved model fits for large galaxies

In previous versions of PHOTO (including that used for the EDR), the exponential and de Vaucouleurs models were fit only to the central 4.4'' radius of each object. This was fine for distant, faint galaxies, but tended to give misleading results for nearby galaxies with large angular extent. The present code now does a much more reasonable fit to large galaxies, meaning that, for example, model magnitudes of bright galaxies are meaningful (but see below). A small bug in the pipeline caused the model fluxes of a handful of bright galaxies to be estimated as negative; this is a rare occurrence, and roughly half the DR1 data have been reduced with a version of the pipeline that fixes this problem.

Model log-likelihoods reported to avoid numerical underflow

These model fits yield a formal likelihood. In previous versions of PHOTO, these had a tendency to underflow to zero, especially for bright objects of high S/N. The current pipeline now outputs the natural log of the likelihood as well, to enable these quantities to be useful.

The "model magnitude bug" and DR1 beta

However, fixing one problem often reveals more subtle, hidden problems. Comparing the model (i.e., exponential and de Vaucouleurs fits) and Petrosian magnitudes of bright galaxies shows a systematic offset of about 0.2 magnitudes (in the sense that the model magnitudes are brighter). This turns out to be due to a bug in the way the PSF was convolved with the models (this bug affected the model magnitudes even when they were fit only to the central 4.4" radius of each object). This caused problems for very small objects (i.e., close to being unresolved). The code forces model and PSF magnitudes of unresolved objects to be the same in the mean by application of an aperture correction, which then gets applied to all objects. The net result is that the model magnitudes are fine for unresolved objects, but systematically offset for galaxies brighter than at least 20th mag.

The bug has been found and fixed in the latest version of the software, but unfortunately, too late for distribution in the DR1 "beta" edition of the Data Archive Server. We have carried out extensive tests with data run through both versions of the code. In summary:

• Colors derived from model magnitudes are almost completely insensitive to the bug. Model magnitudes remain the optimal quantities to use for the colors of extended objects, especially at the faint end.
• For photometry of unresolved sources, one should use PSF magnitudes.
• For photometry of resolved sources, one should use Petrosian magnitudes, especially at the bright end.
• The scale sizes derived from the model fits are systematically wrong. For exponential fits, the effective radii are systematically 0.15" too large in the present code (almost independent of r_e itself), while for the de Vaucouleurs fits, they are roughly 25% too large (almost independent of r_e itself, for r_e > 2 arcsec). These correction factors depend on seeing to some level, of course.

DR1 model magnitudes are better than EDR model magnitudes

Again, it is worth emphasizing that the model magnitudes in DR1 are appreciably better than those used in the EDR. We are reprocessing all the DR1 data with a version of the photometric pipeline that addresses this remaining bug, and will release the results to the public in the final release of DR1 as soon as they are ready. Also see About DR1.

The Astrometric Pipeline

The astrometric pipeline has been improved in a number of ways, including superior treatment of chromatic aberration. That, together with the incorporation of improved centroids from the photometric pipeline and improved astrometric calibration catalogs, have improved the astrometric solutions. Where UCAC data are available, the rms residuals per coordinate in the astrometric solutions are typically of order 60 mas. See Pier et al. (2003) (AJ, or astro-ph/0211375) for a complete description of the astrometric pipeline, including differences in the astrometry between EDR and DR1.

The Spectroscopic Pipelines

Improved spectrophotometry

The 2d spectroscopic pipeline (idlspec2d) has upgraded to much improved flat-fielding, bias subtraction, and handling of bad columns and pixels in the data. Sky subtraction has been improved, especially in the red, by allowing a gradient in the sky brightness across a spectroscopic plate. The spectrophotometric flux calibration is improved as well, as is the correction for absorption lines from the Earth's atmosphere. The spectrophotometry is still not perfect, however; the detailed slope of the continuum of objects can be systematically off by as much as 20% in the flux in the extremes, and one occasionally finds artificial features near the dichroic (i.e., around 6000 Angstroms) of a few percent, and features in the blue as large as 5-10%.

The above statements about the quality of the spectrophotometry hold for point sources. For extended sources, the definition of spectrophotometry itself becomes somewhat ambiguous (see the discussion in Section 4.10 of the EDR paper). In particular, in the presence of spectral gradients, the smear calibrations may be difficult to interpret for extended objects.

Saturated emission lines

On occasion, bright emission lines in objects such as compact star-forming galaxies will saturate, causing line ratios to become meaningless. The relevant pixels are usually (but not always) flagged as having exactly zero flux errors.

Better continuum and line fits

There have been upgrades to the continuum and line-fitting routines in the spectro1d pipeline. Improved stellar templates have much improved the classification of unusual types of stars. Comparison of two parallel versions of the redshift/classification code shows an error rate of only 0.5%, with an additional 1% of spectra of too low signal-to-noise ratio to allow unambiguous classification. The DR1 released spectroscopic data have used this comparison to recognize errors in the redshifts, and update them accordingly; we therefore expect that the number of incorrect redshifts in the final outputs for objects of reasonable signal-to-noise ratio to be no more than a few tenths of a percent.

Spectroscopic classification of galaxies and the eClass parameter

The galaxy spectral classification eigentemplates for DR1 are created from a much larger sample of spectra than were used in the EDR, now approximately 200,000. The eigenspectra used in DR1 are an early version of those created by Yip et al.

eClass is a single-parameter classifier based on the expansion coefficients (eCoeff1-5) (see EDR). The sign of the second eigenspectrum has been reversed with respect to that of EDR; therefore we recommend using the expression atan(-eCoeff2/eCoeff1) rather than eClass as the single-parameter classifier. This will retain the sense of negative values denoting early-type spectra as was the case in EDR and also Yip et al. This will become the definition in spectro rerun 23, which will be contained in the final release of DR1.