% \documentclass[iop,floatfix]{emulateapj}
% \documentclass[iop,floatfix,onecolumn]{emulateapj}
\documentclass[12pt,preprint]{aastex}
% \pdfoutput=1

\RequirePackage{color}
\RequirePackage{code}
\input{astro.sty}

% online version may use color, but print version needs b/w
\def\plotmode{col}
%\def\plotmode{bw}

%\def\plotext{pdf}
\def\plotext{ps}

%\def\picdir{/home/eugene/chipresid.20140404}
\def\picdir{/data/pikake.2/eugene/chipresid.20140404}

% Pick a terse version of the title here;
\shorttitle{PS1 Data Processing System}
\shortauthors{E.A. Magnier et al}
\begin{document}
\title{Pan-STARRS Data Processing System}

% this is a crude trick to get the order of affiliations right.  These
% names are used in the affiliations below.  The user needs to (1) set
% the order and numbers to have the correct sequence in the author
% list and (2) re-order the list at the bottom (and comment-out as needed)
\def\IfA{1}
\def\CfA{2}
\def\MPIA{3}
\def\Princeton{3}
\def\USNO{4}
\def\JHU{1}

% This example has a first author from UH:
\author{
Eugene A. Magnier,\altaffilmark{\IfA}
IPP Team,
%PS Builder List
% W.~S. Burgett,\altaffilmark{\IfA}
% K.~C. Chambers,\altaffilmark{\IfA} 
% L. Denneau,\altaffilmark{\IfA}
% P. Draper,\altaffilmark{\DUR}
% H.~A. Flewelling,\altaffilmark{\IfA}
% T. Grav,\altaffilmark{\IfA}
% J. N. Heasley,\altaffilmark{\IfA}
% K. W. Hodapp,\altaffilmark{\IfA}
% M. E. Huber,\altaffilmark{\IfA}
% R. Jedicke,\altaffilmark{\IfA}
% N. Kaiser,\altaffilmark{\IfA}
% R.-P. Kudritzki,\altaffilmark{\IfA}
% G. A. Luppino,\altaffilmark{\IfA}
% R. H. Lupton,\altaffilmark{\Princeton}
% E. A. Magnier,\altaffilmark{\IfA}
% N. Metcalfe,\altaffilmark{\DUH}
% D. G. Monet,\altaffilmark{\USNO}
% J.~S. Morgan,\altaffilmark{\IfA}
% P. M. Onaka,\altaffilmark{\IfA}
% P.~A. Price,\altaffilmark{\Princeton}
% C.~W. Stubbs,\altaffilmark{\CfA}
% W.~E. Sweeney,\altaffilmark{\IfA}
% J.~L. Tonry, \altaffilmark{\IfA}
% R. J. Wainscoat,\altaffilmark{\IfA} and 
% C. Z. Waters,\altaffilmark{\IfA}
} % this bracket terminates author list

% The ordering here should be sequential, matching the sequence in the list of authors:
\altaffiltext{\IfA}{Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu HI 96822}
% \altaffiltext{\CfA}{Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138}
% \altaffiltext{\Princeton}{Department of Astrophysical Sciences, Princeton University, Princeton, NJ 08544, USA}
% \altaffiltext{\USNO}{US Naval Observatory, Flagstaff Station, Flagstaff, AZ 86001, USA}
% \altaffiltext{\JHU}{Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA}
% \altaffiltext{\MPIA}{Max Planck Institute for Astronomy, K\"onigstuhl 17, D-69117 Heidelberg, Germany}
\begin{abstract}

Lorem ipsum dolor sit amet, consectetur adipiscing elit. Vestibulum
bibendum nisi id tristique posuere. Duis eu mollis nulla. Maecenas est
turpis, mattis tempor urna vitae, placerat rhoncus sem. Lorem ipsum
dolor sit amet, consectetur adipiscing elit. Sed quis velit
nisl. Aliquam erat volutpat. Cras lacinia, nisl tristique auctor
molestie, dolor nulla rhoncus purus, ac accumsan nunc nunc ac
nibh. Maecenas vitae mollis mauris. Ut sollicitudin pulvinar purus,
eget luctus lorem tincidunt vitae. Vestibulum eu mattis neque. Nulla
in tortor id urna dapibus gravida a vel leo.

\end{abstract}

% insert additional keywords as appropriate:
\keywords{Surveys:\PSONE }

% \section{INTRODUCTION}\label{sec:intro}

\section{IPP Software Subsystems}

\subsection{Processing Database}

A critical element in the Pan-STARRS IPP infrastructure is the
processing database.  This database, using the mysql database engine,
tracks information about each of the processing stages.  It is used as
the point of mediation of all IPP operations.  Processing stages
within the IPP perform queries of the database to identify the data to
be processed at a given stage.  As the processing for a particular
stage is completed, summary information about the stage is written
back to the database.  In this way, the database records this history
of the processing, and also provides the information needed to
successive processing stages to begin their own tasks.  

The processing database is colloquially referred to as the `gpc1'
database, since a single instance of the database is used to track the
processing of images and data products related to the PS1 GPC1 camera.
This same database engine also has instances for other cameras
processed by the IPP, e.g., GPC2, the test cameras TC1, TC3, the
Imaging Sky Probe (ISP), etc.

Within the processing database, the various processing stages are
represented as a set of tables.  In general, there is a top level
table which defines the conceptual list of processing items either to
be done, in progress, or completed.  An associated table lists the
details of elements which have been processed.  For example, one
critical stage is the Chip processing stage, discussed below, in which
the individual chips from an exposure are detrended and sources are
detected.  Within the gpc1 database, there is a top-level table called
`chipRun' in which each exposure has a single entry.  Associated with
this table is the `chipProcessedImfile' table, which contains one row
for each of the (up to 60) chips associated with the exposure.  The
top-level tables, such as chipRun, are populated once the system has
decided that a specific item (e.g., an exposure) should be processed
at that stage.  Initially, the entry is given a state of `run',
denoting that the exposure is ready to be processed.  The low-level
table entries, such as the chipProcessedImfile entries, are only
populated once the element (e.g., the chip) has been processed by the
analysis system.  Once all elements for a given stage, e.g., chips in
this case, are completed, then the status of the top-level table entry
(chipRun) are switched from `run' to `done'.

If the analysis of an element (e.g., chip) completed successfully,
then the corresponding table row (e.g., chipProcessedImfile) is listed
with a fault of 0.  If the analysis failed, then a non-zero fault is
recorded.  An analysis which results in a fault is one in which the
failure is thought to be temporary.  For example, if a computer had a
network glitch and was unable to write out some of the output files,
this would be an ephemeral failure which was not a failing of the
data, but merely the processing system.  On the other hand, if the
analysis failed because of a problem with the input data, this is
noted by setting a non-zero value in a different table field,
`quality'.  For example, if the chip analysis failed to discover any
stars because the image was completely saturated, the analysis can
complete successfully (fault = 0), but the `quality' field will be set
to a non-zero value.  The various processing stages are able to select
only the good (quality = 0) elements of a prior stage when choosing
items for processing.  For example, the Camera calibration stage will
only use data from chips with good quality data, dropping the failed
chips from the rest of the analysis.  On the other hand, a fault in
one of the elements for a given stage will block any dependent stages
from processing that item.  In this way, if a glitch occurs and a chip
from an exposure failed to be written to disk in the Chip stage, the
system will not partially process the exposure with the rest of the
chips.  Since many of the faults which occur are ephemeral, the
processing stages are set up to occasional clear and re-try the
faulted entries.  Thus, automatic processing is able to keep the data
flowing even in the face of occasional network glitches or hardware
crashes.

\subsection{Nebulous}

\subsection{Pantasks \& Parallel Processing}

\subsection{DVO}

The Pan-STARRS IPP uses an internal database system, distinct from the
publically visible database system, to determine the association
between multiple detections of the same astronomical object and as
part of the astrometric and photometric calibration process.  This
database system, called the ``Desktop Virtual Observatory'' (DVO) was
developed originally for the LONEOS project, and used as part of the
CFHT Elixir system (Magnier \& Cuillandre REF).  The capabilities of
this databasing system have been somewhat expanded for the Pan-STARRS
context.  

One of the main purposes of the DVO system is to define the
relationship between individual detections of an astronomical object
and the definition of that object.  Before describing the database
schema, we will discuss the process by which detections are associated
with objects.  New detections are generally added to the database in a
group associated with, for example, an image from the GPC1 camera.  As
new detections are loaded, they are compared to the objects already
stored in the database.  If an object is already found in the database
within the match radius, the new detection is associated to that
object. If more than one object exists within the database, the
detection is associated with the closest object.  

Detections in DVO have a special piece of metadata called the
\code{photcode} which identifies the source of the measurement.  A
\code{photcode} has a name which in general consists of the name of
the camera or telescope (e.g., GPC1 or 2MASS), the name (or short-hand
name) of the filter used for the measurement (e.g., $g$), and an
identifier for the detector, if not unique (e.g., XY01 for GPC1).
Along with each name, there is a numerical value for the photcode.  A
table within the DVO system, \code{Photcode}, lists the photcoes and
defines a number of additional pieces of information for each
photcode.  These include the nominal zero point and airmass slope, as
well as color trends to transform a measurement in the specific
photcode to a common system.  There are 3 classes of photcodes defined
within the DVO system.  Those photcodes associated with detections
from an image loaded into the database system are called \code{DEP}
photcodes.  There are also photcodes associated with the average
photometry values, called SEC photcodes.  There are also those
measurements which come from external data sources for which DVO does
not have any information to determine a calibration (e.g.,
instrumental magnitudes and detector coordinates).  These are
measurements are reference values and are assigned REF photcodes.

In the implementation of DVO used for the PV3 calibration analysis,
the database tables are stored on disk using binary FITS tables.  Each
type of database table is stored as a separate file, or a collection
of files for table which are spatially partitioned.  The binary FITS
tables may be optionally compressed using the (to date) experimental
FITS binary table compression strategy outlined by REF.  In this
compression scheme, using a strategy similar to that used for FITS
image compression (REF), the data stored in the binary table is
compressed and stored in the 'HEAP' section of the FITS table.  In
brief, each column in the FITS table is compressed as one (or more)
blocks.  The standard fields which describe the data column format
(e.g., TFORM1) are replaced with columns which describe the location
and size of the compressed data in the HEAP section; the information
about the uncompressed data is moved to a field with 'Z' prepended
(e.g., ZFORM1) and an additional field is added to define the
compression algorithm (e.g., ZCTYP1).  The column names (e.g., TTYPE1)
and units (e.g., TUNIT1) are retained in their original form.  The
compression algorithm can treat the entire column as a single block of
data, or it may be broken into a number of chunks, each compressed in
turn (this must be the same for all columns).  Additional header
information is added to describe the block sizes and infomation needed
to describe the HEAP data section.  The compression algorithms
currently defined consist of the GZIP, RICE, PLIO, and HCOMPRESS
(REFS).  For GZIP, the compression algorithm may transpose the byte
order before compression: for floating point data of a similiar
dynamic range, this choice may allow for better compression as each
byte in the 4 or 8 byte floating point value is more likely to be
similar to the same byte in other rows than to the other bytes of the
same row value.  This option is called \code{GZIP_2} in the FITS
standard, as opposed to the standard order, \code{GZIP_1}.  The DVO
system can be set to specify the compression options for each column
in the tables.  In practice, we have chosen a default in which
floating point numbers used \code{GZIP_2}, character strings use
\code{GZIP_1}, integers use \code{RICE}.  

\subsubsection{Sky Partition}

DVO includes two major classes of database tables: those containing
information directly about astronomical objects in the sky and those
containing other supporting information.  The object-related tables
are partitioned on the basis of position in the sky: objects within a
region bounded by lines of constant RA,DEC are contained in a specific
file.  The boundaries and the associated partition names are stored in
one of the supporting tables, \code{SkyTable}.  This table contains
the definitions of the boundaries for each sky region (\code{R_MIN},
\code{R_MAX}, \code{D_MIN}, \code{D_MAX}), the name of the sky region,
an ID (\code{INDEX}, equal to the sequence number of the region in the
table), and index entries to enable navigation within the table.  The
regions are defined in a hierarchical sense, with a series of levels
each containing a finer mesh of regions covering the sky.  

In the default used by the PV3 DVO, the partitioning scheme is based
on the one used by the Hubble Space Telescope Guide Star Catalog
files.  Level 0 is a single region covering the full sky.  Level 1
divides the sky in Declination into bands 7.5\degree\ high.  Level 2
subdivides these Declination bands in the RA direction, with spacing
related to the stellar density.  Level 3 divides these RA chunks into
4 - 8 smaller partitions.  This level exactly matches the HST GSC
layout, and uses the same naming convention to identify the
partitions: n0000/0000, etc. \note{more on the names?}.  Level 4
further divides these regions by a factor of 16.  In the
\code{SkyTable}, a region at one level has a pointer to its parent
region (the one which contains it) and a sequence pointing to its
children (regions it contains).  The \code{SkyTable} enables fast
lookups of the on-disk partitions which map to a specific coordinate
on the sky.  In general, a single DVO will have the full sky
represented with tables at a single level, though it is possible for
mixed levels to be used, this mode is not well tested.  For the PV3
master database, the partitioning at the 5th level results in \approx
150,000 regions to cover the full sky, of which \approx 110,000 are
used for the PV3 $3\pi$ data.  The densest portions of the bulge
contain at most \approx 300k astronomical objects in the database
files, with an associated maximum of 30M measurements in these files.
With the compression scheme described above, this makes the largest
database files \approx 3GB, which can be loaded into memory in 30
seconds on our partition machines.

The DVO software system allows the tables which are partitioned across
the sky to also be distributed across multiple computers, which we
call partition hosts.  A single file defines the names of these
partition hosts and the location of the database partition on the
disks of that machine.  The \code{SkyTable} contains elements to
define by ID the parition host to which a partitioned set of tables
has been assigned.  Operations which query the database, or perform
other operations on the database, are aware of the partitioning scheme
and will launch their operations as remote processes on the machines
which contain the data they need.  For example, a query for data from
a small region will launch sub-query operations on the machines which
contain the data overlapping the region of interest.  These remote
query operations will select the database information which matches
the query request (i.e., applying restrictions as defined) and return
to the master process the results.  The results from the various
partition hosts are then merged into a single result by the master
process.  This parallelization is critical to querying and
manipulating the enormous database on a reasonable timescale.

\subsection{Tables which describe objects} 

Two tables carry the most important information about the astronomical
objects in the database: Average and SecFilt.  These two tables
specify the main average properties of the astronomical object.  The
Average table includes the astrometric information ($\alpha, \delta,
\mu \alpha, \mu \delta, \pi$) and associated errors, data quality
flags for each object, links to the other tables, and a number of IDs,
with one row for each astronomical object.  \note{go into complete
  detail here on the IDs?}.  The SecFilt table\footnote{The name
  SecFilt is a bit of a historical misnomer: originally, DVO was
  designed for a monochromatic survey and data for a single
  photometric band was maintained in the Average table.  Later, DVO
  was adapted to a multifilter system and additional filters were
  added to the SecFilt (Secondary Filter) table.  Eventually, the
  schema was normalized and all photometric data placed in SecFilt,
  with the Primary filter concept being dropped, but the name has
  since stuck.} contains average photometric information for a
collection of filters.  A given DVO instance has a specified set of
filters for which average photometry is stored in the SecFilt table.
The number and choice of filters for the SecFilt may be modified by
the database administrator fairly easily, but the process of updating
the database is somewhat expensive (\approx 24 hours for the current
PV3 instance).  Thus the choice is semi-static for a given DVO
instance.  In the case of the PV3 DVO instance, 9 average bandpasses
are defined: {\it g, r, i, z, y, J, H, K, w}.  The SecFilt table
contains one row for each filter for each object, thus the PV3 DVO
contains 9 times as many rows as the Average table.  The order of the
table is fixed in relation to the Average table: row $i$ of Average
defines the object with photometry contained in rows $9i \rightarrow 9i +
8$ ($i$ is zero counting).  

The individual measurements of the astronomical objects are carried in
the table \code{Measure}.  This table lists the values measured by
\code{psphot} for each chip, warp, or stack image.  This includes the
instrumental magnitudes for the PSF, aperture, and Kron photometry;
raw position (Xccd, Yccd) and second moments (Mxx, Myy, Mxy), along
with shape parameters of the PSF model at the position of the object
(FWx, FWy, theta).  This table also includes metadata such as the
exposure time, the date \& time of the observation, airmass, azimuth,
and information specifying the filter \note{describe the photcodes}.
The \code{Measure} table also carried the calibration magnitude offsts
($M_{\rm cal}$ and $M_{\rm flat}$ discussed below) and the
astrometrically calibrated position.  Astrometric offsets for several
systematic corrections discussed below are also defined for each
measurement.  Since stacks and forced warp photometry may have
non-significant values, the table is somewhat de-normalized in that it
also carried instrumental flux values for the PSF, aperture, and Kron
photometry.  

In the \code{Measure} table, there are three fields which provide two
independent links from the specific measurement to the associated
object: \code{Measure.catID} specifies the spatial partition to which
the measurement belongs; \code{Measure.objID} specifies to which entry
in the \code{Average} table the measurement belongs.  These two 32 bit
fields can thus be combined into a single 64 bit unique ID for all
objects in the database.  In addition, the field \code{Measure.averef}
specifies the row number in the \code{Average} table of the associated
object.  The \code{Measure} table may be unsorted, in which case it is
slow to find the measurements associated with a specific object (a
full table scan is required).  After the table is sorted, the
\code{Measure} rows for a given object are grouped together.  In the
case, the fields \code{Average.measureOffset} and
\code{Average.Nmeasure} define an index for the code to jump to the
list of measurements for a single object.  The field
\code{Measure.imageID} defines the link from the measurement to the
image which supplied the measurement.

\note{some discussion of the db construction, addstar, dvomerge, etc?}

For the warp images, we also measure the weak lensing KSB parameters
related to the shear and smear tensors.  These measurements are stored
in the \code{Lensing} table, along with the radial aperture fluxes for
radii numbers 5, 6, \& 7 (XX, XX, XX arcsec).  This table contains one
row for every warp row.  Similarly to the \code{Measure} table, the fields
\code{objID}, \code{catID}, and \code{averef} define links from the
\code{Lensing} table to the \code{Average} table.  In a similar
fashion, the fields \code{Average.lensingOffset} and
\code{Average.Nlensing} are the index into the sorted \code{Lensing}
table entries.  \note{discuss failure of the Lensing to Measure
  indexing}

The values stored in the \code{Lensing} table are used to calculate
average values for each of these types of measurements in each
filter.  The \code{Lensobj} table stores the averaged KSB and radial
aperture photometry for each of the 5 filters \grizy.  This table
contains one entry per object per filter.  The table is not generally
stored unsorted as it is calculated after the full database is
populated.  The link from \code{Average} to \code{Lensobj} is defined
by the fields \code{Average.offsetLensobj} and
\code{Average.Nlensobj}.  Each \code{Lensobj} row also includes the
photcode (filter) for which the average lensing (and radial aperture)
properties have been calculated. 

The \code{Galphot} table stores the results of the forced galaxy
fitting analysis for each object that has been measured.  This table
contains one row per filter and model type (Sersic, Exponential,
DeVaucouleur) if forced galaxy models have been calculate for the
object.  \note{need to expand on this somewhat}

The \code{Starpar} table carries measurements provide by Greg Green \&
Eddie Schlafly from their analysis of the SED of objects in the PS1
$3\pi$ data, using the \note{PV1?} version of the analysis (Green et
al REF).  In this work, the goal was a 3D model of the dust in the
Galaxy based on Pan-STARRS (\note{and WISE \& 2MASS?}) photometry.  As
part of this analysis, the authors fit the SEDs of all \note{stellar?}
sources with stellar models including free parameters of extinction,
distance modulus, metallicity, and absolute r-band magnitude.  While
these photometric distance modulus measurements are not extremely
precise (see below), they provide a constraint on the distance is used
in our analysis of the astrometry (see Section~\ref{sec:astrometry}).

\subsection{Other Tables} 

Data from GPC1 (and other cameras processed by the IPP) are loaded
into DVO in units \code{smf} files generated by the Camera calibration
stage.  As described above, these files contain all measurements from
a complete exposure, with measurements from each chip grouped into
separate FITS tables.  When these measurements are loaded into the
\code{Measure} and similar tables, a subset of the information from
the chip header is used to populated a row in the DVO \code{Image}
table.  This table contains one row for each chip known to DVO, with
information such as the filter (\code{photcode}), the exposure time,
the airmass, the astrometric calibration terms, the photometric
zero point, etc.  For GPC1 and other mosaic cameras, an additional row
is defined to carry the projection and camera distortion elements of
the astrometry model.  As chips are loaded into this table, they are
assigned an internal ID (a running sequence in the table).  Images may
also be assigned an external ID: in the case of the GPC1 images, this
ID is defined by the processing mysql database and is guaranteed to be
unique within the processing system. 

Other tables are used to track information used by the calibration
system.  This includes the complete set of flat-field corrections
determined by the photometry calibration analysis (see
Section~\ref{sec:relphot}) and the astrometric flat-field corrections
determined by the astrometry calibration analysis (see Section~\ref{sec:relastro})

\section{IPP Data Processing Stages}

\subsection{Download from Summit}

As exposures are taken by the PS1 telescope \& camera system, the 60
OTA CCDs are read out by the camera software system and each chip is
written to disk on computers at the summit in the PS1 facility.  The
chip images are written to a collection of machines in the PS1
facility called the `pixel servers'.  After the images are written to
disk, a summary listing of the information about the exposure and the
chip images are written to an http server system called the
`datastore'.  The datastore exposes, via http, a list of the exposures
obtained since the start of the PS1 operations.  Requests to this
server may restrict to the latest by time.  Each row in the listing
includes basic information about the exposure: an exposure identifier
(e.g., o5432g0123o; see~\ref{GPC1.names} for details), the date and
time of the exposure, \note{etc}.  The row also provides a link to a
listing of the chips associated with that exposure.  This listing
includes a link to the individual chip FITS files as well as an md5
CHECKSUM.  Systems which are allowed access may download chip FITS
files via http requests to the provided links.

During night-time operations, while the telescope is observing the sky
and the camera subsystem is saving images to the pixel servers and
adding their information to the datastore list, the IPP subsystem
called `summitcopy' monitors the datastore in order to discover new
exposures ready for download.  Once a new exposure has been listed on
the datastore, summitcopy adds an entry of the exposure to a table in
the processing database (summitExp).  This tells the summitcopy to
look for the list of chips, which are then added to another table
(summitImfile).  The summitcopy system then attempts to download the
chips from the summit pixel servers with an http request.  As the chip
files are downloaded, their md5 checksum values are calculated and
compared with the value reported by the camera subsystem / datastore.
Download failures are rare and marked as a non-zero fault, allowing for a
manual recovery, rather than automatically rejecting the failed
chips.  

\subsection{Image Registration}

Once chips for an exposure have all been downloaded, the exposure is
ready to be registered.  In this context, `registration' refers to the
process of adding them to the database listing of known, raw exposures
(not to be confused with 'registration' in the sense of pixel
re-alignment).  The result of the registration analysis is an entry
for each exposure in the rawExp table, and one for each chip in the
rawImfile table.  These tables are critical for downstream processing
to identify what exposures are available for processing in any other
stage.  In the registration stage, a large amount of descriptive
metadata for each chip is added to the rawImfile table, some of which
is extracted from the chip FITS file headers (e.g., RA, DEC, FILTER)
and some of which is determined by a quick analysis of the pixels
(e.g., mean pixel values, standard deviation).  The chip-level
information is merged into a set of exposure-level metadata and added
to the rawExp table entry.  The exposure-level metadata may be the
same as any one of the chip, in a case where the values are duplicated
across the chip files (e.g., the name of the telescope or the date \&
time of the exposure), or it may be a calculation based on the values
from each chip (e.g., average of the average pixel values).

Unlike much of the rest of the IPP stage, the raw exposures may only
have a single entry in the registration tables of the processing
database tables (rawExp and rawImfile).

\subsection{Chip Processing}

The science analysis of an exposure begins with the processing of the
individual chips, the Chip Processing stage.  This analysis step has
two main goals: the removal of the instrumental signature from the
pixel values (detrending) and the detection of the sources in the
objects.  In the Chip stage, the individual chips are processed
independently in parallel within the data processing cluster.  Within
the processing computer cluster, most of the data storage resources
are in the form of computers with large raids as well as substantial
processing capability.  The processing system attempts to locate one
copy of specific raw chips on pre-defined computers for each chip.
The processing system is aware of this data localization and attempts
to target the processing of a particular chip to the machine on which
the data for that chip is stored.  The output products are then
primarily saved back to the same machine.  This `targetted' processing
was an early design choice to minimize the system wide network load
during processing.  In practice, as computer disks filled up at
different rates, the data has not been localized to a very high
degree.  The targeted processing has probably reduced the network load
somewhat but it has not been as critical of a requirement as
originally expected.

The Chip processing stage consists of: reading the raw image into
memory, appyling the detrending steps (see~\note{Waters et al}),
stiching the individual OTA cells into a single chip image, detection
and characterization of the sources in the image (see~\note{Magnier et
  al}), and output of the various data products.  These include the
detrended chip image, variance image, and mask image, as well as the FITS
catalog of detected sources.  The PSF model and background model are
also saved, along with a processing log.  A selection of summary
metadata describing the processing results are saved and written to
the processing database along with the completion status of the
process.  Finally, binned chip images are generated (on two scales,
binned by 16 and 256 pixels) for use in the visualization system of
the processing monitor tool.

\subsection{Camera Calibration}

After sources have been detected and measured for each of the chip,
the next stage is to perform a basic calibration of the full exposure.
This stage starts with the collection of FITS tables containing the
instrumental measurements of the detected sources, primarily the
positions ($x_{\rm ccd}, y_{\rm ccd}$) and the instrumental PSF
magnitudes.  The data for all chips of an exposure are loaded by the
analysis program.  The header information is used to determine the
coordinates of the telescope boresite (RA, DEC, Position angle).
These three coordinates are used, along with a model of the camera
layout, to generate an initial guess for the astrometry of each chip.
Reference star coordinates and magnitudes are loaded from a reference
catalog for a region corresponding to the boundaries of the exposure,
padded by a large fraction of the exposure diameter in case of a
modest pointing error.  The guess astrometry is used to match the
reference catalog to the observed stellar positions in the focal plane
coordinate system.  Once an acceptable match is found, the astrometric
calibration of the individual chips is performed, including a a fit to
a single model for the distortion introduced by the camera optics.
After the astrometic analysis is completed, the photometric
calibration is determined using the final match to the reference
catalog.  At this stage, pre-determined color terms may be included to
convert the reference photometry to an appropriate photometric
system.  For PS1, this is used to generate synthetic w-band photometry
for areas where no PS1-based calibrated w-band photometry is
available.  For more details, see \note{Magnier et al}.

In addition to the astrometric and photometric calibrations, the
Camera stage also generates the dynamic masks for the images.  The dynamic
masks include masking for optical ghosts, glints, saturated stars,
diffraction spikes, and electronic crosstalk.  The mask images
generated by the Chip stage are updated with these dynamic masks and a
new set of files are saved for the downstream analysis stages.

The Camera stage also merges the binned chip images
(see~\ref{sec:chip}) into single jpeg images of the entire focal
plane.  These jpeg images can then be displayed by the process
monitoring system to visualize the data processing.

\subsection{Warp}

Once astrometric and photometric calibrations have been performed,
images are geometrically transformed into a set of common pixel-grid
images with simple projections from the sky.  These images, called
skycells, can then be used in subsequent stacking and difference image
analysis without concern about the astrometric transformation of an
exposure.  This processing is called `warping'; the warp analysis
stage is run on all exposures before they are processed further.  For
details on the warping algorithm, see \note{Waters et al paper}.

The output products from the Warp stage consist of the skycell images
containing the signal, the variance, and the mask information.  These
images have been shipped to STScI and \note{are available / will be
  available} from the image extraction tools \note{in DR2}.

\subsection{Stack}

The skycell images generated by the Warp process are added together to
make deeper, higher signal-to-noise images in the Stack stage.  The
stacks also fill in coverage gaps between different exposures,
resulting in an image of the sky with more uniform coverage than a
single exposure.  See~\note{Waters paper} for details on the stack
combination algorithm.

In the IPP processing, stacks may be made with various options for the
input images.  During nightly science processing, the 8 exposures per
filter for each Medium Deep field are combined into a set of stacks
for that field.  These so-called `nightly stacks' are used by the
transient survey projects to detect the faint supernovae, among other
transient events.  For the PV3 $3\pi$ analysis, all filter images from
the $3\pi$ survey observation were stacked together to generate a
single set of images with $\sim 10 - 20\times$ the exposure of the
individual survey exposures.  The signal, variance, and mask images
resulting from these deep stacks are part of the DR1 release and are
available from the image extraction tools.

For the PV3 processing of the Medium Deep fields, stacks have been
generated for the nightly groups and for the full depth using all
exposures (deep stacks).  In addition, a 'best seeing' set of stack
have been produced \note{using image quality cuts to be described}.
We have also generated out-of-season stacks for the Medium Deep
fields, in which all image not from a particular observing season for
a field are combined into a stack.  These later stacks are useful as
deep templates when studying long-term transient events in the Medium
Deep fields as they are not (or less) contaminated by the flux of the
transients from a given season.

\subsection{Stack Photometry}

The stack images are generated in the Stack stage of the IPP, but the
source detection and extraction analysis of those images is deferred
until a separate stage, the Stack Photometry stage.  This separation
is maintained because the stack photometry analysis is performed on
all 5 filter stack images at the same time.  By deferring the
analysis, the processing system may decouple the generation of the
pixels from the source detection.  This makes the sequencing of
analysis somewhat easier and less subject to blocks due to a failure
in the stacking analysis.

The stack photometry algorithms are described in detail in
\note{Magnier et al}.  In short, sources are detected in all 5 filter
images down to the $5\sigma$ significance.  The collection of detected
sources is merged into a single master list.  If a source is detected
in at least two bands, or only in $y$-band, then a PSF model is fitted
to the pixels of the other bands in which the source was not detected.
This forced photometry results in lower significance measurements of
the flux at the positions of objects which are thought to be real
sources, by virtue of triggering a detection in at least two bands.
The relaxed limit for $y$-band is included to allow for searches of
$y$-dropout objects: it is known that faint, high-redshift quasars may
be detected in $y$-band only.  The casual user of the PV3 dataset
should be wary of sources detected only in $y$-band as these are
likely to have a higher false-positive rate than the other stack
sources.

The stack photometry output files consist of a set of FITS tables in a
single file, with one file for each filter.  Within one of these
files, the tables include: the measurements of sources based on the
PSF model; aperture like parameters such as the Petrosian flux and
radius; the convolved Galaxy model fits; the radial aperture
measurements.  \note{is this list complete?}

The stack photometry output catalogs are re-calibrated for both
photometry and astrometry in a process very similar to the Camera
calibration stage.  In the case of the stack calibration, however,
each skycell is processed independently.  The same reference catalog
is used for the Camera and Stack calibration stages.

\subsection{Forced Warp Photometry}

Traditionally, projects which use multiple exposures to increase the
depth and sensitivity of the observations have generated something
equivalent to the stack images produced by the IPP analysis.  In
theory, the photometry of the stack images produces the `best'
photometry catalog, with best sensitivity and the best data quality at
all magnitudes (c.f, CFHT Legacy survey, COSMOS, etc).  In practice,
the stack images have some significant limitations due to the
difficulty of modelling the PSF variations.  This difficulty is
particularly severe for the Pan-STARRS $3\pi$ survey stacks due to the
combination of the substantial mask fraction of the individual
exposures, the large instrinsic image quality variations within a
single exposure, and the wide range of image quality conditions under
which data were obtained and used to generate the $3\pi$ PV3 stacks.

For any specific stack, the point spread function at a particular
location is the result of the combination of the point spread
functions for those individual exposures which went into the stack at
that point.  Because of the high mask fraction, the exposures which
contributed to pixels at one location may be somewhat different just a
few tens of pixels away.  Because of the intrinsic variations in the
PSF across an exposure and because of the variations from exposure to
exposure, the distribution of point spread functions of the images
used at one position may be quite different from those at a nearby
location.  In the end, the stack images have a effective point spread
function which is not just variable, but changing significantly on
small scales in a highly textured fashion.  

Any measurement which relies on a good knowledge of the PSF at the
location of an object either needs to determine the PSF variations
present in the stack, or the measurement will be somewhat degraded.
The highly textured PSF variations make this a very challenging
problem: not would such a PSF model require an unusually fine-grained
PSF model, there would likely not be enough PSF stars in an given
stack to determine the model at the resolution required.  The IPP
photometry analysis code uses a PSF model with 2D variations using a
grid of at most $6\times 6$ samples per skycell, a number reasonably
well-matched to the density of stars at most moderate Galactic
latitudes.  This scale is far too large to track the fine-grained
changes apparent in the stack images.

Thus PSF photometry as well as convolved Galaxy models in the stack
are degraded by the PSF variations.  Aperture-like measurements are in
general not as affected by the PSF variations, as long as the aperture
in question is large compared to the FWHM of the PSF.

%% The IPP team initially explored the option of convolving each input
%% warp to a single target PSF chosen to match the worst of the input
%% images for a given stack.  

The PV3 $3\pi$ analysis solves this problem by using the sources
detected in the Stack images and performing forced photometry on the
individual warp images used to generate the stack.  This analysis is
performed on all warps for a single filter as a single job, though
this is more of a bookkeepping aid as it is not necessary for the
analysis of the different warps to know about the results of the other
warps.

In the forced warp photometry, the positions of sources are loaded
from the stack outputs.  PSF stars are pre-identified and a PSF model
generated for each warp based on those stars, using the same stars for
all warps to the extent possible (PSF stars which are excessively
masked on a particular image are not used to model the PSF).  The PSF
model is fitted to all of the known source positions in the warp
images.  Aperture magnitudes, Kron magnitudes, and moments are also
measured at this stage for each warp.  Note that the flux measurement
for a faint, but significant, source from the stack image may be at a
low significance ($< 5\sigma$) in any individual warp image; the flux
may even be negative for specific warps.  When combined together,
these low-significance measurements will result in a signficant
measurement as the signal-to-noise increases by $\sqrt{N}$.  

\subsection{Forced Galaxy Models}

The convolved galaxy models are also re-measured on the warp images by
the forced photometry analysis stage.  In this analysis, the galaxy
models determined by the stack photometry analysis are used to seed
the analysis in the individual warps.  The purpose of this analysis is
the same as the forced PSF photometry: the PSF of the stack is poorly
determined due to the masking and PSF variations in the inputs.
Without a good PSF model, the PSF-convolved galaxy models are of
limited accuracy.  

In the forced galaxy model analysis, we assume that the galaxy
position and position angle, along with the Sersic index if
appropriate, have been sufficiently well determined in the stack
analysis.  In this case, the goal is to determine the best values for
the major and minor axis of the elliptical contour and at the same
time the best normalization corresponding to the best elliptical shape
(and thus the best galaxy magnitude value).

For each warp image, the Stack value for the major and minor axis are
used as the center of a $7\times 7$ grid search of the major and minor
axis parameter values.  The grid spacing is defined as a function of
the signal-to-noise of the galaxy in the stack image so that bright
galaxies are measured with a much finer grid spacing that faint
galaxies \note{need to quantify this}.  For each grid point, the major
and minor axis values at that point are determined for the model.  The
model is then generated and convolved with the PSF model for the warp
image at that point.  The resulting model is then compared to the warp
pixel data values and the best fit normalization value is defined.
The normalization and the $\chi^2$ value for each grid point is
recorded.  

For a given galaxy, the result is a collection of $\chi^2$ values for
each of the grid points spanning all warp images.  A single $\chi^2$
grid can then be made from all warps by combining each grid point
across the warps.  The combined $\chi^2$ for a single grid point is
simply the sum of all $\chi^2$ values at that point.  If, for a single
warp image, the galaxy model is excessively masked, then that image
will be dropped for all grid points for that galaxy.  The reduced
$\chi^2$ values can be determined by tracking the total number of warp
pixels used across all warps to generate the combined $\chi^2$ values.
From the combined grid of $\chi^2$ values, the point in the grid with
the minimum $\chi^2$ is found.  Quadratic interpolation is used to
determine the major, minor axis values for the interpolated minimum
$\chi^2$ value.  The errors on these two parameters is then found by
determining the contour at which the \note{reduced?} $\chi^2$
increases by 1.  

Thus the Forced Galaxy Model analysis uses the PSF information from
each warp to determine a best set of convovled galaxy models for each
object in the stack images.  \note{discuss the subset of galaxy models
  and objects}.

\subsection{Difference Images}

Two of the primary science drivers for the Pan-STARRS system are the
search hazardous asteroids and the search for Type Ia supernovae to
measure the history of the expansion of the universe.  Both of these
projects require the discovery of faint, transient source in the
images.  For the hazardous asteroids, and solar system studies in
general, the sources are transient because they are moving between
observations; supernovae are stationary but transient in brightness.
In both cases, the discovery of these sources can be enhanced by
subtracting a static reference image from the image taken at a certain
epoch.  The quality of such a difference image can be enhanced by
convolving one or both of the images so that the PSFs in the two
images are matched.  \note{discuss Alard-Lupton}. 

In the Difference Image stage, the IPP generates diffferece images for
specified pairs of images.  It is possible for the difference image to
be generated from a pair of warp images, from a warp and a stack of
some variety, or from a pair of stacks.  During the PS1 survey, pairs
of exposures, call TTI pairs (see~\note{Survey Strategy}), were
obtained for each pointing within a $\approx$ 1 hour period in the
same filter, and to the extent possible with the same orientation and
boresite position.  The standard PS1 nightly processing generated
difference images from the resulting warp pairs (`warp-warp diffs').

The nightly stacks generated for the Medium Deep fields were combined
with a template reference stack image to generate `stack-stack diffs'
for these fields each night.  

For the PV3 processing, the entire collection of warps for the $3\pi$
survey were combined with the $3\pi$ stacks to generate `warp-stack
diffs'.  

\subsection{Addstar : DVO Ingest}

\subsection{Calibration Operations}

\subsection{IPP to PSPS}

\subsection{PSPS Load \& Merge}

\section{IPP Hardware Systems}

\subsection{Kihei Processing Cluster} 

\subsection{Los Alamos National Labs} 

\subsection{UH Cray Cluster} 

\end{document}
