Changeset 39868
- Timestamp:
- Dec 15, 2016, 3:48:40 PM (10 years ago)
- Location:
- trunk/doc/release.2015
- Files:
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- 8 edited
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Makefile.Common (modified) (2 diffs)
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inputs/lib.bib (modified) (2 diffs)
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ps1.analysis/Makefile (modified) (1 diff)
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ps1.analysis/analysis.tex (modified) (32 diffs)
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ps1.calibration/Makefile (modified) (2 diffs)
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ps1.calibration/calibration.tex (modified) (41 diffs)
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ps1.datasystem/Makefile (modified) (1 diff)
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ps1.datasystem/datasystem.tex (modified) (5 diffs)
Legend:
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trunk/doc/release.2015/Makefile.Common
r39866 r39868 7 7 PS2PDF_OPTS = "-dAutoFilterColorImages=false -dColorImageFilter=/FlateEncode" 8 8 9 ifeq ($(DO_PDFLATEX),1) 10 MY_LATEX = $(PDFLATEX) 11 else 12 MY_LATEX = $(PSLATEX) 13 endif 14 9 15 %.pdf: %.tex 10 $(PSLATEX) $*.tex 11 $(BIBTEX) $* 12 $(PSLATEX) $*.tex 13 # thumbpdf --modes=dvips $*.pdf 14 # $(PSLATEX) $*.tex 15 dvips -z -t letter -o $*.ps $*.dvi 16 ps2pdf $(PS2PSF_OPT) $*.ps $*.pdf 16 $(MY_LATEX) $*.tex 17 if [ $(DO_BIBTEX) -eq 1 ]; then $(BIBTEX) $*; fi 18 $(MY_LATEX) $*.tex 19 if [ $(DO_BIBTEX) -eq 1 ]; then $(MY_LATEX) $*.tex; fi 20 thumbpdf --modes=dvips $*.pdf 21 $(MY_LATEX) $*.tex 22 if [ $(DO_PDFLATEX) -eq 0 ]; then dvips -z -t letter -o $*.ps $*.dvi; fi 23 if [ $(DO_PDFLATEX) -eq 0 ]; then ps2pdf $(PS2PDF_OPTS) $*.ps $*.pdf; fi 17 24 # @rm -f $*.ps $*.dvi $*.aux $*.log $*.tbr $*.tbd $*.toc $*.tpm $*.lof body.tmp head.tmp 18 25 26 %.tgz: 27 tar --transform 's%inputs/%%' -zcf $@ $(FILES) 19 28 clean : 20 29 $(RM) *.bib *.log *.dvi *.aux *.toc *.tbd *.tbr *.tpm *.lof *.out *~ core body.tmp head.tmp … … 24 33 25 34 empty: clean 35 36 foo: 37 @echo all: $^ 38 @echo 1st: $< 39 @echo file: $@ 40 @echo word: $* -
trunk/doc/release.2015/inputs/lib.bib
r39866 r39868 9987 9987 } 9988 9988 9989 % moffat profile: 9990 @ARTICLE{1969A&A.....3..455M, 9991 author = {{Moffat}, A.~F.~J.}, 9992 title = "{A Theoretical Investigation of Focal Stellar Images in the Photographic Emulsion and Application to Photographic Photometry}", 9993 journal = {\aap}, 9994 year = 1969, 9995 month = dec, 9996 volume = 3, 9997 pages = {455}, 9998 adsurl = {http://adsabs.harvard.edu/abs/1969A%26A.....3..455M}, 9999 adsnote = {Provided by the SAO/NASA Astrophysics Data System} 10000 } 10001 10002 % more moffat example: 10003 @ARTICLE{1983A&A...126..278B, 10004 author = {{Buonanno}, R. and {Buscema}, G. and {Corsi}, C.~E. and {Ferraro}, I. and 10005 {Iannicola}, G.}, 10006 title = "{Automated photographic photometry of stars in globular clusters}", 10007 journal = {\aap}, 10008 keywords = {Astronomical Photography, Globular Clusters, Star Distribution, Stellar Spectrophotometry, Hertzsprung-Russell Diagram, Stellar Evolution}, 10009 year = 1983, 10010 month = oct, 10011 volume = 126, 10012 pages = {278-282}, 10013 adsurl = {http://adsabs.harvard.edu/abs/1983A%26A...126..278B}, 10014 adsnote = {Provided by the SAO/NASA Astrophysics Data System} 10015 } 10016 10017 % daophot 9989 10018 @ARTICLE{1987PASP...99..191S, 9990 10019 author = {{Stetson}, P.~B.}, … … 15800 15829 } 15801 15830 15831 % dophot 15802 15832 @ARTICLE{1993PASP..105.1342S, 15803 15833 author = {{Schechter}, P.~L. and {Mateo}, M. and {Saha}, A.}, -
trunk/doc/release.2015/ps1.analysis/Makefile
r39848 r39868 1 1 # $Id: Makefile,v 1.16 2006-01-16 01:11:40 eugene Exp $ 2 3 DO_PDFLATEX = 0 4 DO_BIBTEX = 1 2 5 3 6 help: 4 7 @echo "USAGE: make (target)" 5 @echo " targets: all analysis"8 @echo " targets: all pdf tgz" 6 9 7 all: analysis.pdf 8 analysis: analysis.pdf 10 all: pdf tgz 11 pdf: analysis.pdf 12 tgz: analysis.tgz 9 13 10 ANALYSIS = analysis.tex 14 FILES = \ 15 ../inputs/astro.sty \ 16 ../inputs/code.sty \ 17 ../inputs/apj.bst \ 18 ../inputs/lib.bib \ 19 peaks.ps \ 20 FWHM.smooth.trend.ps1.ps \ 21 moment.class.ps \ 22 radial.profiles.ps \ 23 analysis.tex \ 24 analysis.bbl 11 25 12 # pics/Metadata.ps 13 # pics/earthrot.ps 14 15 analysis.pdf: $(ANALYSIS) 16 17 analysis.ps: $(ANALYSIS) 26 analysis.pdf: $(FILES) 27 analysis.tgz: $(FILES) 18 28 19 29 include ../Makefile.Common -
trunk/doc/release.2015/ps1.analysis/analysis.tex
r39866 r39868 1 1 \documentclass[iop,floatfix]{emulateapj} 2 % \documentclass[iop,floatfix,onecolumn]{emulateapj} 2 3 3 % \pdfoutput=1 4 4 5 % see latex.readme.txt for notes on using the PS1 template6 %\documentclass[12pt,preprint]{aastex}7 %\documentclass[manuscript]{aastex}8 %\documentclass[preprint2]{aastex}9 %\documentclass[preprint2,longabstract]{aastex}10 5 \RequirePackage{color} 11 6 \RequirePackage{code} … … 18 13 %\def\plotext{pdf} 19 14 \def\plotext{ps} 15 \def\plottype{eps} 20 16 21 17 %\def\picdir{/home/eugene/chipresid.20140404} … … 33 29 % list and (2) re-order the list at the bottom (and comment-out as needed) 34 30 \def\IfA{1} 31 \def\Princeton{2} 32 \def\DUR{3} 35 33 \def\CfA{2} 36 \def\MPIA{3}37 \def\Princeton{2}38 \def\USNO{4}39 \def\JHU{1}40 34 41 35 % This example has a first author from UH: 42 36 \author{ 43 37 Eugene A. Magnier,\altaffilmark{\IfA} 44 R. H. Lupton,\altaffilmark{\Princeton}38 % R. H. Lupton,\altaffilmark{\Princeton} 45 39 W.~E. Sweeney,\altaffilmark{\IfA} 46 40 K.~C. Chambers,\altaffilmark{\IfA} … … 49 43 P.~A. Price,\altaffilmark{\Princeton} 50 44 C. Z. Waters,\altaffilmark{\IfA} 51 PS1 Builders 45 % PS1 Builders 46 L. Denneau,\altaffilmark{\IfA} 47 P. Draper,\altaffilmark{\DUR} 48 R. Jedicke,\altaffilmark{\IfA} 49 K. W. Hodapp,\altaffilmark{\IfA} 50 R.-P. Kudritzki,\altaffilmark{\IfA} 51 N. Metcalfe,\altaffilmark{\DUR} 52 C.~W. Stubbs,\altaffilmark{\CfA} 52 53 % W.~S. Burgett,\altaffilmark{\IfA} 53 % K.~C. Chambers,\altaffilmark{\IfA}54 % L. Denneau,\altaffilmark{\IfA}55 % P. Draper,\altaffilmark{\DUR}56 % H.~A. Flewelling,\altaffilmark{\IfA}57 54 % T. Grav,\altaffilmark{\IfA} 58 55 % J. N. Heasley,\altaffilmark{\IfA} 59 % K. W. Hodapp,\altaffilmark{\IfA}60 % M. E. Huber,\altaffilmark{\IfA}61 % R. Jedicke,\altaffilmark{\IfA}62 56 % N. Kaiser,\altaffilmark{\IfA} 63 % R.-P. Kudritzki,\altaffilmark{\IfA}64 57 % G. A. Luppino,\altaffilmark{\IfA} 65 58 % R. H. Lupton,\altaffilmark{\Princeton} 66 59 % E. A. Magnier,\altaffilmark{\IfA} 67 % N. Metcalfe,\altaffilmark{\DUH}68 60 % D. G. Monet,\altaffilmark{\USNO} 69 61 % J.~S. Morgan,\altaffilmark{\IfA} 70 62 % P. M. Onaka,\altaffilmark{\IfA} 71 % P.~A. Price,\altaffilmark{\Princeton}72 % C.~W. Stubbs,\altaffilmark{\CfA}73 % W.~E. Sweeney,\altaffilmark{\IfA}74 63 % J.~L. Tonry, \altaffilmark{\IfA} 75 % R. J. Wainscoat,\altaffilmark{\IfA} and 76 % C. Z. Waters,\altaffilmark{\IfA} 64 R. J. Wainscoat\altaffilmark{\IfA} 77 65 } % this bracket terminates author list 78 66 79 67 % The ordering here should be sequential, matching the sequence in the list of authors: 80 68 \altaffiltext{\IfA}{Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu HI 96822} 81 % \altaffiltext{\CfA}{Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138}82 69 \altaffiltext{\Princeton}{Department of Astrophysical Sciences, Princeton University, Princeton, NJ 08544, USA} 70 \altaffiltext{\DUR}{Department of Physics, Durham University, South Road, Durham DH1 3LE, UK} 71 \altaffiltext{\CfA}{Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138} 83 72 % \altaffiltext{\USNO}{US Naval Observatory, Flagstaff Station, Flagstaff, AZ 86001, USA} 84 73 % \altaffiltext{\JHU}{Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA} … … 157 146 described in detail in \cite{2012ApJ...750...99T}. 158 147 148 {\color{red} {\em Note: These papers are being placed on arXiv.org to 149 provide crucial support information at the time of the public 150 release of Data Release 1 (DR1). We expect the arXiv versions to 151 be updated prior to submission to the Astrophysical Journal in 152 January 2017. Feedback and suggestions for additional information 153 from early users of the data products are welcome during the 154 submission and refereeing process.}} 155 156 \section{Background} 157 159 158 The photometric and astrometric precision goals for the Pan-STARRS\,1 160 159 surveys were quite stringent: photmetric accuracy of 10 … … 183 182 astrometry. 184 183 185 \subsection{Comparable Programs}186 187 184 A variety of astronomical software packages perform the basic object 188 185 detection, measurement, and classification tasks needed by the … … 198 195 pro: well-tested, stable code. con: limited range of models, 199 196 algorithm converges slowly to a PSF model, limited tests of PSF 200 validity, inflexible code base, fortran (P. Schechter)197 validity, inflexible code base, fortran \citep{1993PASP..105.1342S}. 201 198 202 199 \item DAOPhot : Pixel-map PSF model with analytical component. pro: 203 200 well-tested, high-quality photometry. con: Difficult to use in an 204 automated fashion, does it handle 2D variations well? (P. Stetson)201 automated fashion, does it handle 2D variations well? \citep{1987PASP...99..191S}. 205 202 206 203 \item Sextractor : pure aperture measurement with rudimentary object 207 204 subtraction. pro: fast, widely used, easy to automate. con: poor 208 205 object separation in crowded regions, PSF-modeling was only in beta, 209 not widely used at the time. (E. Bertin) 210 211 \item apphot : IRAF-based aperture photometry. pro: widely used. 212 con: IRAF-based, aperture photometry. (???) 206 not widely used at the time \citep{sextractor}. 213 207 214 208 \item galfit : detailed galaxy modeling. not a multi-object PSF 215 209 analysis tool. con: does not provide a PSF model, not easily 216 210 automated. very detailed results in very slow processing. only a 217 galaxy analysis program . (C. Impey)211 galaxy analysis program \citep{2002AJ....124..266P}. 218 212 219 213 \item SDSS phot : con: tightly integrated into the SDSS software 220 environment . (R. Lupton)214 environment \citep{2001ASPC..238..269L}. 221 215 222 216 \end{itemize} … … 425 419 subtracted might be useful for detection or even analysis of brighter 426 420 sources. Table~\ref{tab:mask_values} lists the 16 bit values used for 427 PS1 mask images, along with their description (see \note{Waters et428 a l. paper} for additional information).421 PS1 mask images, along with their description \citep[see][for 422 additional information]{waters2017}. 429 423 430 424 \begin{table*} … … 495 489 which the values of \code{SKY} and \code{SKY_SIGMA} are calculated for 496 490 each object in the output catalog. See also the discussion in 497 \ note{Waters et al REF}.491 \cite{waters2017}. 498 492 499 493 \subsection{Initial Object Detection} … … 580 574 \begin{figure}[htbp] 581 575 \begin{center} 582 \includegraphics[width=\hsize,angle=0,clip]{peaks.ps} 583 \caption{Illustration of peak finding and culling peaks within a 576 % \includegraphics[type=\plottype,ext=.\plotext,width=3.5in,height=2.5in,viewport=60 60 560 310]{peaks} 577 % \includegraphics[type=\plottype,ext=.\plotext,width=\hsize,height=1in,viewport=60 60 560 310,clip]{peaks} 578 \includegraphics[type=\plottype,ext=.\plotext,width=\hsize,height=0.5\hsize,viewport=60 60 560 310,clip]{peaks} 579 \caption{\label{fig:peaks} Illustration of peak finding and culling peaks within a 584 580 footprint. Insignificant peaks within the footprint of a brighter 585 581 peak are ignored in further processing. } … … 608 604 \code{FOOTPRINT_CULL_NSIGMA_DELTA} (4.0) sigmas below the peak of 609 605 interest, the peak is considered to be {\em locally insignificant} and 610 removed from the list of possible detections. In the vicinity of a 611 saturated star, the rule is somewhat more agressive as the flat-topped 612 or structured saturated top of a bright star may appear as multiple 613 peaks with highly significant cols between them. However, this is an 614 artifact of the proximity to saturation. In this regime, we require 615 the col to also be a fixed fraction (5\%) of the saturation below the 616 peak to avoid being marked as locally insignificant. 606 removed from the list of possible detections (see 607 Figure~\ref{fig:peaks}). In the vicinity of a saturated star, the 608 rule is somewhat more agressive as the flat-topped or structured 609 saturated top of a bright star may appear as multiple peaks with 610 highly significant cols between them. However, this is an artifact of 611 the proximity to saturation. In this regime, we require the col to 612 also be a fixed fraction (5\%) of the saturation below the peak to 613 avoid being marked as locally insignificant. 617 614 618 615 \subsubsection{Centroid and higher-order Moments} 616 \label{sec:moments} 619 617 620 618 \begin{figure}[htbp] 621 619 \begin{center} 622 \includegraphics[ width=\hsize,angle=0,clip]{FWHM.smooth.trend.ps1.ps}623 \caption{ Example of the biases encountered when measuring the second620 \includegraphics[type=\plottype,ext=.\plotext,width=\hsize,height=2.0\hsize,viewport=60 60 413 760]{FWHM.smooth.trend.ps1} 621 \caption{\label{fig:moments.window} Example of the biases encountered when measuring the second 624 622 moments. A simulated image was generated using the PS1 PSF 625 623 profile. Each panel corresponds to a different value of … … 656 654 signal-to-noise of the object. 657 655 658 These effects are illustrated in Figure~\ref{fig:moment .window} using656 These effects are illustrated in Figure~\ref{fig:moments.window} using 659 657 simulated data. An image was generated with a PSF model matching the 660 658 radial profile of the PS1 PSF model with a FWHM of 1.4 arcseconds. As … … 736 734 these moments. 737 735 738 The Kron radius is defined the be 2.5$\times$ the first radial moment. 739 The Kron flux is the sum of (sky-subtracted) pixel fluxes within the 740 Kron radius. We also calculate the flux in two related annular 741 apertures: the Kron inner flux is the sum of pixel values for the 742 annulus $R_1 < r < 2.5 R_1$, while the Kron outer flux is the sum of 743 pixel values for $2.5 R_1 < r < 4 R_1$. The first radial moment is 744 limited at the low and high ends by $R_{\rm min} < M_r < R_{\rm max}$ 745 where $R_{\rm min}$ is the first radial moment of the PSF stars, or 746 0.75$\times$ \code{MOMENTS_GAUSS_SIGMA} if that cannot be 747 determined. $R_{\rm max}$ is set to \code{PSF_MOMENTS_RADIUS}, the 748 size of the moments aperture. 736 The Kron radius \citep{1980ApJS...43..305K} is defined the be 737 2.5$\times$ the first radial moment. The Kron flux is the sum of 738 (sky-subtracted) pixel fluxes within the Kron radius. We also 739 calculate the flux in two related annular apertures: the Kron inner 740 flux is the sum of pixel values for the annulus $R_1 < r < 2.5 R_1$, 741 while the Kron outer flux is the sum of pixel values for $2.5 R_1 < r 742 < 4 R_1$. The first radial moment is limited at the low and high ends 743 by $R_{\rm min} < M_r < R_{\rm max}$ where $R_{\rm min}$ is the first 744 radial moment of the PSF stars, or 0.75$\times$ 745 \code{MOMENTS_GAUSS_SIGMA} if that cannot be determined. $R_{\rm 746 max}$ is set to \code{PSF_MOMENTS_RADIUS}, the size of the moments 747 aperture. 749 748 750 749 \subsection{PSF Determination} … … 906 905 \begin{figure}[htbp] 907 906 \begin{center} 908 \includegraphics[ width=\hsize,angle=0,clip]{moment.class.ps}907 \includegraphics[type=\plottype,ext=.\plotext,width=\hsize,height=\hsize,viewport=60 60 560 560]{moment.class} 909 908 \caption{\label{fig:moment.class} Illustration of PSF star selection using the FWHM derived 910 909 from the second moments in $X_{\rm ccd}$ and $Y_{\rm ccd}$ … … 919 918 \subsubsection{PSF Candidate Object Model Fits} 920 919 920 % \note{link to psLibADD} 921 921 922 All candidate PSF objects are then fitted with the selected object 922 923 model, allowing all of the parameters (PSF and independent) to vary in 923 924 the fit. PSPhot uses the Levenberg-Marquardt minimization technique 924 \note{link to psLibADD}for the non-linear fitting. Non-linear925 for the non-linear fitting. Non-linear 925 926 fitting can be very computationally intensive, particularly for if the 926 927 starting parameters are far from the minimization values. PSPhot uses … … 1012 1013 \subsubsection{PSF Model applied to detected objects} 1013 1014 1014 \note{review the discussion below}1015 % \note{review the discussion below} 1015 1016 1016 1017 Once a PSF model has been selected for an image, PSPhot attempts to 1017 1018 fit all of the detected objects, above a user-defined signal-to-noise 1018 ratio (\note{KEYWORD}) with the PSF model. For these fits, the1019 dependent parameters are fixed by the PSF model and only the 4 1020 independent object model parameters are allowed to vary in the fit. 1021 PSPhot again uses Levenberg-Marquardt minimization for the non-linear 1022 fitting. The objects are fitted in their S/N order, starting with the 1023 brightest andworking down to the user-specified limit.1019 ratio with the PSF model. For these fits, the dependent parameters 1020 are fixed by the PSF model and only the 4 independent object model 1021 parameters are allowed to vary in the fit. PSPhot again uses 1022 Levenberg-Marquardt minimization for the non-linear fitting. The 1023 objects are fitted in their S/N order, starting with the brightest and 1024 working down to the user-specified limit. 1024 1025 1025 1026 Once a solution has been achieved for an object, PSPhot attempts to … … 1108 1109 1109 1110 \subsubsection{Source Size Assessment} 1111 \label{sec:source.size} 1110 1112 1111 1113 After the PSF model has been fitted to all sources, and the Kron flux … … 1294 1296 \frac{y^2}{2\sigma_y^2} + \sigma_{\rm xy} x y $). The Pseudo-Gaussian 1295 1297 is a Taylor expansion of the Gaussian and is used by Dophot 1296 \citep{dophot}. The latter profiles are similar to the Moffat profile 1297 form \citep{moffat,buonanno}, with small differences. For the PS1 1298 GPC1 analysis, we used the \code{PS1_V1} model, which we found by 1299 experimentation to match well to the observed profiles generated by 1300 PS1. Using a fixed power-law exponent results in somewhat faster 1301 profile fitting compared to the variable power-law exponent model. 1298 \citep{1993PASP..105.1342S}. The latter profiles are similar to the 1299 Moffat profile form \citep{1969AA.....3..455M,1983AA...126..278B}, 1300 with small differences. For the PS1 GPC1 analysis, we used the 1301 \code{PS1_V1} model, which we found by experimentation to match well 1302 to the observed profiles generated by PS1. 1303 Figure~\ref{fig:radial.profiles} shows example radial profiles for 1304 moderately bright stars in fairly good (0.9 arcsec) and poor (2.2 1305 arcsec) seeing. Using a fixed power-law exponent results in somewhat 1306 faster profile fitting compared to the variable power-law exponent 1307 model. 1302 1308 1303 1309 % moffat : 1969A&A.....3..455M … … 1306 1312 \begin{figure}[htbp] 1307 1313 \begin{center} 1308 \includegraphics[ width=\hsize,angle=0,clip]{radial.profiles.ps}1309 \caption{ Radial profiles of stellar images from PS1. These two1314 \includegraphics[type=\plottype,ext=.\plotext,width=\hsize,height=\hsize,viewport=60 60 560 560]{radial.profiles} 1315 \caption{\label{fig:radial.profiles} Radial profiles of stellar images from PS1. These two 1310 1316 profiles illustrate the radial trend of the PS1 PSFs for a star 1311 1317 with FWHM 0.9 arcsec (red) and 2.2 arcsec (blue). The black line … … 1372 1378 \code{RMAX_NN}). 1373 1379 1374 \note{these profiles are not saved in PSPS}1380 % \note{these profiles are not saved in PSPS} 1375 1381 1376 1382 \subsection{Petrosian Radii and Magnitudes} 1377 1383 1378 Petrosian (REF) defined an adaptive aperture based on a ratio of 1379 surface brightnesses. The motivation is to define an aperture which 1380 can be determined for galaxies without significant biases as a 1381 function of distance. Since surface brightness in a resolved object 1382 is conserved, using a ratio of surface brightness to define a spatial 1383 scale results in a spatial scale which is constant regardless of 1384 galaxy distance. 1384 \cite{1976ApJ...209L...1P} defined an adaptive aperture based on a 1385 ratio of surface brightnesses. The motivation is to define an 1386 aperture which can be determined for galaxies without significant 1387 biases as a function of distance. Since surface brightness in a 1388 resolved object is conserved, using a ratio of surface brightness to 1389 define a spatial scale results in a spatial scale which is constant 1390 regardless of galaxy distance. 1385 1391 1386 1392 To measure the Petrosian radius and flux, we start by defining a … … 1428 1434 median) flux in the annulus is within 1 $\sigma$ of the local sky 1429 1435 level. If this limit is not reached before the slope of the flux from 1430 one annulus to the next is less tha t \note{SOMETHING,1431 psphotRadialProfileWings.c}, then the annulus at which the slope1432 reaches this limit is used to define the sky radius. These values are 1433 saved in the output smf / cmf files, but not sent to the PSPS. The1434 sky radius value is used below in thecalculation of the kron magnitude.1436 one annulus to the next is less than a user-defined limit, then the 1437 annulus at which the slope reaches this limit is used to define the 1438 sky radius. These values are saved in the output smf / cmf files, but 1439 not sent to the PSPS. The sky radius value is used below in the 1440 calculation of the kron magnitude. 1435 1441 1436 1442 \subsection{Kron Magnitudes} 1437 1443 1438 Preliminary Kron radius and flux values are calculated soon after1439 sources are detected (\ref{sec:moments}). However, these preliminary 1440 values are not accurate due to the window-functions applied. After 1441 sources have been characterized and the PSF model is well-determined, 1442 the Kron parameters are re-calculated more carefully. In this version1443 of the calculation, the image is first smoothed by Gaussian kernel 1444 with $\sigma = 1.7$ pixels, corresponding to a FWHM of 1.0\arcsec in 1445 the PS1 stack images. Next, the Kron radius is determined in an 1446 iterative process: the first radial moment is measured using the pixels in an 1447 aperture 6$\times$ the first radial moment from the previous 1448 iteration. On the first iteration, the sky radius is used in place of 1449 the first radial moment. By default, 2 iterations are performed. The 1450 Kron radius is defined the be 2.5$\times$ the first radial moment. 1451 The Kron flux is the sum of pixel fluxes within the Kron radius. We 1452 also calculate the flux in two related annular apertures: the Kron1453 inner flux is the sum of pixel values for the annulus $R_1 < r < 2.5 1454 R_1$, while the Kron outer flux is the sum of pixel values for $2.5 1455 R_1 < r < 4 R_1$. 1444 Preliminary Kron radius and flux values \citep{1980ApJS...43..305K} 1445 are calculated soon after sources are detected (Section~\ref{sec:moments}). 1446 However, these preliminary values are not accurate due to the 1447 window-functions applied. After sources have been characterized and 1448 the PSF model is well-determined, the Kron parameters are 1449 re-calculated more carefully. In this version of the calculation, the 1450 image is first smoothed by Gaussian kernel with $\sigma = 1.7$ pixels, 1451 corresponding to a FWHM of 1.0\arcsec\ in the PS1 stack images. Next, 1452 the Kron radius is determined in an iterative process: the first 1453 radial moment is measured using the pixels in an aperture 6$\times$ 1454 the first radial moment from the previous iteration. On the first 1455 iteration, the sky radius is used in place of the first radial moment. 1456 By default, 2 iterations are performed. The Kron radius is defined 1457 the be 2.5$\times$ the first radial moment. The Kron flux is the sum 1458 of pixel fluxes within the Kron radius. We also calculate the flux in 1459 two related annular apertures: the Kron inner flux is the sum of pixel 1460 values for the annulus $R_1 < r < 2.5 R_1$, while the Kron outer flux 1461 is the sum of pixel values for $2.5 R_1 < r < 4 R_1$. 1456 1462 1457 1463 Two details in the calculation above should be noted. First, for … … 1460 1466 calculations. The window used for the calculations is constrained to 1461 1467 be at least the size of the aperture based on the PSF stars 1462 ( \ref{sec:moments}). At the other extreme, noise may make the radius1468 (Section~\ref{sec:moments}). At the other extreme, noise may make the radius 1463 1469 grow excessively, resulting in an unrealistically low effective 1464 1470 surface brightness. The aperture is constrained to be less than a … … 1471 1477 opposites sides of the central pixel are considered together. The 1472 1478 geometric mean of the two fluxes is used to replace the flux values. 1473 If the object has 180\degree symmetry, this operation has no impact.1479 If the object has 180\degree\ symmetry, this operation has no impact. 1474 1480 However, if one of the two pixels is unusually high, the value will be 1475 1481 surpressed by the matched pixel on the other side. This trick has the … … 1480 1486 1481 1487 In the galaxy model fittting stage, sources which meet certain 1482 criteria are fitted with analytical models for galaxies. The 1483 three models used for the PV3 analysis have similar form: 1488 criteria are fitted with analytical models for galaxies. Three 1489 traditional analytical galaxy models are implemented in \code{psphot} 1490 and used in the PV3 analysis: 1484 1491 \begin{itemize} 1485 1492 \item Exponential profile : $f = I_0 e^{-\rho}$ 1486 \item DeVaucouleur profile : $f = I_0 e^{-\rho^{1/4}}$1487 \item Sersic : $f = I_0 e^{-\rho^{1/n}}$1493 \item DeVaucouleur profile \citep{1948AnAp...11..247D}: $f = I_0 e^{-\rho^{1/4}}$ 1494 \item Sersic \citep{1963BAAA....6...41S} : $f = I_0 e^{-\rho^{1/n}}$ 1488 1495 \end{itemize} 1489 1496 where $\rho$ is a normalized radial term: $\rho = … … 1500 1507 our best guess for the PSF model at the location of the galaxy. For 1501 1508 the PV3 analysis, all sources detected in the 'bright source' analysis 1502 step ($S/N > 20 ?$) were fitted with all three galaxy models, unless 1503 (a) the morphological test identified the source as a likely cosmic 1504 ray (\ref{CR}) or (b) the peak of the PSF profile was above the 1505 saturation limit for the chip \note{(link to the handling of 1506 saturation in detrend paper)}. Sources in the denser portions of 1507 the Galactic plane and bulge were not included in the analysis. This 1508 restriction limited the total time spent on the galaxy modeling 1509 analysis at the expense of galaxy photometry in the plane (though Kron 1510 photometry is available for those objects). The Galactic Plane region 1511 was defined by $|b| > b_{\rm min}$ where $b_{\rm min} = b_0 + r_b 1512 e^{\frac{-l^2}{2 \sigma_b^2}}$. For the PV3 analysis, $b_0 = XX$, 1513 $r_b = XX$, $\sigma_b = XX$. 1509 step ($S/N > 20$) were fitted with all three galaxy models, unless (a) 1510 the morphological test identified the source as a likely cosmic ray 1511 (Section~\ref{sec:source.size}) or (b) the peak of the PSF profile was 1512 above the saturation limit for the chip \citep[see the discussion in 1513 ][ regarding the masking of saturated pixels]{waters2017}. Sources in 1514 the denser portions of the Galactic plane and bulge were not included 1515 in the analysis. This restriction limited the total time spent on the 1516 galaxy modeling analysis at the expense of galaxy photometry in the 1517 plane (though Kron photometry is available for those objects). The 1518 Galactic Plane region was defined by $|b| > b_{\rm min}$ where $b_{\rm 1519 min} = b_0 + r_b e^{\frac{-l^2}{2 \sigma_b^2}}$. For the PV3 1520 analysis, $b_0 = $20\degree, $r_b = $15\degree, $\sigma_b = $50\degree. 1521 1522 % \note{need a discussion of the detector saturation behavior 1523 1524 % \note{more detail below?} 1514 1525 1515 1526 Before the non-linear fitting may be performed, it is necessary to … … 1521 1532 ($R_{xx}$, $R_{yy}$ , $R_{xy}$) values; it was found that such a guess 1522 1533 tended to be too small and resulted in more iterations rather than 1523 fewer. \note{more detail on that?}The 1st radial moment (see1534 fewer. The 1st radial moment (see 1524 1535 \ref{sec:moments}) is used to estimate the effective radius of the 1525 1536 model based on the results of Graham \& Driver (2005, Table 1). They … … 1606 1617 For the small size of the PSF model used to convolve the galaxy model 1607 1618 images, it was found that this direct convolution was faster than 1608 using an FFT-based convolution \note{(examples?)} 1619 using an FFT-based convolution. 1620 1621 % \note{(examples?)} 1609 1622 1610 1623 For the Exponential and DeVaucouleur fits, all parameters are fitted … … 1656 1669 for all 5 filters. In this analysis, the best model for each object 1657 1670 is subtracted from the image pixels for all objects excluding the 1658 object in consideration. The 'best model' is \note{TBD}. 1671 object in consideration. The 'best model' is determined based on the 1672 minimum $\chi^2$ value for the model fits. 1673 1674 % \note{more discussion of the selection of the best model}. 1659 1675 1660 1676 In addition to the raw radial apertures, the stack images are each … … 1667 1683 procedure is then repeated with a target FWHM of 8\arcsec. 1668 1684 1669 \note{is the first convolution done with the Alard-Lupton technique?} 1685 % \note{is the first convolution done with the Alard-Lupton technique?} 1686 1687 \acknowledgments 1688 1689 The Pan-STARRS1 Surveys (PS1) have been made possible through 1690 contributions of the Institute for Astronomy, the University of 1691 Hawaii, the Pan-STARRS Project Office, the Max-Planck Society and its 1692 participating institutes, the Max Planck Institute for Astronomy, 1693 Heidelberg and the Max Planck Institute for Extraterrestrial Physics, 1694 Garching, The Johns Hopkins University, Durham University, the 1695 University of Edinburgh, Queen's University Belfast, the 1696 Harvard-Smithsonian Center for Astrophysics, the Las Cumbres 1697 Observatory Global Telescope Network Incorporated, the National 1698 Central University of Taiwan, the Space Telescope Science Institute, 1699 the National Aeronautics and Space Administration under Grant 1700 No. NNX08AR22G issued through the Planetary Science Division of the 1701 NASA Science Mission Directorate, the National Science Foundation 1702 under Grant No. AST-1238877, the University of Maryland, and Eotvos 1703 Lorand University (ELTE) and the Los Alamos National Laboratory. 1704 1705 \bibliographystyle{apj} 1706 % \bibliography{lib}{} 1707 \input{analysis.bbl} 1708 1709 \end{document} 1670 1710 1671 1711 \subsection{Forced Photometry : PSFs} … … 1675 1715 \subsection{Output Options} 1676 1716 1677 \note{need to discuss tests}1678 1679 \note{need to discuss failings and holes}1717 % \note{need to discuss tests} 1718 1719 % \note{need to discuss failings and holes} 1680 1720 1681 1721 \section{Alternative Scenarios} … … 1759 1799 \end{verbatim} 1760 1800 1761 \bibliographystyle{apj}1762 \bibliography{lib}{}1763 1764 \end{document}1765 1766 1801 Figures Needed for this document: 1767 1802 … … 1791 1826 * put engineering docs (psLib, psModules) on public website 1792 1827 1793 % example of 2 image figure:1794 \begin{figure}1795 \centering1796 \begin{minipage}{0.45\hsize}1797 \includegraphics[width=0.9\hsize,angle=0,clip]{images/o5677g0123o_XY11_bt_trail.png}1798 \end{minipage}%1799 \begin{minipage}{0.45\hsize}1800 \includegraphics[width=0.9\hsize,angle=0,clip]{images/o5677g0124o_XY11_bt_trail.png}1801 \end{minipage}1802 \caption{Example of a profile cut along the y-axis through a bright star on exposure o5677g0123o OTA11 in cell xy60 (left panel) and on the subsequent exposure o5677g0124o (right panel). In both figures, the green points show the image corrected with all appropriate detrending steps, but without burntool applied, illustrating the amplitude of the persistence trails. The red points show the same data after the burntool correction, which reduces the impact of these features. Both exposures are in the \gps{} filter with exposure times of 43s}1803 \end{figure}1804 -
trunk/doc/release.2015/ps1.calibration/Makefile
r39841 r39868 1 1 # $Id: Makefile,v 1.16 2006-01-16 01:11:40 eugene Exp $ 2 3 DO_PDFLATEX = 0 4 DO_BIBTEX = 1 2 5 3 6 help: … … 5 8 @echo " targets: all calibration" 6 9 7 all: calibration.pdf 8 calibration: calibration.pdf 10 all: pdf tgz 11 pdf: calibration.pdf 12 tgz: calibration.tgz 9 13 10 CALIBRATION = calibration.tex 14 FILES = \ 15 ../inputs/astro.sty \ 16 ../inputs/code.sty \ 17 ../inputs/apj.bst \ 18 ../inputs/lib.bib \ 19 calibration.tex \ 20 calibration.bbl 11 21 12 # pics/Metadata.ps 13 # pics/earthrot.ps 14 15 calibration.pdf: $(CALIBRATION) 16 17 calibration.ps: $(CALIBRATION) 22 calibration.pdf: $(FILES) 23 calibration.tgz: $(FILES) 18 24 19 25 include ../Makefile.Common -
trunk/doc/release.2015/ps1.calibration/calibration.tex
r39865 r39868 1 %\documentclass[iop,floatfix]{emulateapj}1 \documentclass[iop,floatfix]{emulateapj} 2 2 % \pdfoutput=1 3 3 4 4 % see latex.readme.txt for notes on using the PS1 template 5 \documentclass[12pt,preprint]{aastex}5 %\documentclass[12pt,preprint]{aastex} 6 6 %\documentclass[manuscript]{aastex} 7 7 %\documentclass[preprint2]{aastex} … … 32 32 % list and (2) re-order the list at the bottom (and comment-out as needed) 33 33 \def\IfA{1} 34 \def\CfA{2} 35 \def\MPIA{3} 36 \def\Princeton{3} 37 \def\USNO{4} 38 \def\JHU{1} 34 \def\LBL{2} 35 \def\Hubble{3} 36 \def\ITC{4} 37 \def\Harvard{5} 38 \def\MPIA{6} 39 \def\ARI{7} 40 \def\Princeton{8} 41 \def\DUR{9} 42 \def\CfA{10} 39 43 40 44 % This example has a first author from UH: 41 45 \author{ 42 Eugene A. Magnier,\altaffilmark{\IfA} 43 IPP Team, 44 %PS Builder List 46 Eugene. A. Magnier,\altaffilmark{\IfA} 47 Edward. F. Schlafly,\altaffilmark{\LBL,\Hubble} 48 Douglas P. Finkbeiner,\altaffilmark{\ITC,\Harvard} 49 J.~L. Tonry,\altaffilmark{\IfA} 50 B. Goldman,\altaffilmark{\MPIA} 51 S. R\"oser,\altaffilmark{\ARI} 52 E. Schilbach,\altaffilmark{\ARI} 53 K.~C. Chambers,\altaffilmark{\IfA} 54 H.~A. Flewelling,\altaffilmark{\IfA} 55 M. E. Huber,\altaffilmark{\IfA} 56 P.~A. Price,\altaffilmark{\Princeton} 57 W.~E. Sweeney,\altaffilmark{\IfA} 58 C. Z. Waters,\altaffilmark{\IfA} 59 % PS1 Builders 60 L. Denneau,\altaffilmark{\IfA} 61 P. Draper,\altaffilmark{\DUR} 62 K. W. Hodapp,\altaffilmark{\IfA} 63 R. Jedicke,\altaffilmark{\IfA} 64 R.-P. Kudritzki,\altaffilmark{\IfA} 65 N. Metcalfe,\altaffilmark{\DUR} 66 C.~W. Stubbs,\altaffilmark{\CfA} 45 67 % W.~S. Burgett,\altaffilmark{\IfA} 46 % K.~C. Chambers,\altaffilmark{\IfA}47 68 % T. Grav,\altaffilmark{\IfA} 48 69 % J. N. Heasley,\altaffilmark{\IfA} 49 % K. W. Hodapp,\altaffilmark{\IfA}50 % R. Jedicke,\altaffilmark{\IfA}51 % H.~A. Flewelling,\altaffilmark{\IfA}52 70 % N. Kaiser,\altaffilmark{\IfA} 53 % R.-P. Kudritzki,\altaffilmark{\IfA}54 71 % G. A. Luppino,\altaffilmark{\IfA} 55 72 % R. H. Lupton,\altaffilmark{\Princeton} … … 57 74 % J.~S. Morgan,\altaffilmark{\IfA} 58 75 % P. M. Onaka,\altaffilmark{\IfA} 59 % P.~A. Price,\altaffilmark{\Princeton} 60 % W.~E. Sweeney,\altaffilmark{\IfA} 61 % C.~W. Stubbs,\altaffilmark{\CfA} 62 % J.~L. Tonry, \altaffilmark{\IfA} 63 % R. J. Wainscoat,\altaffilmark{\IfA} and 76 R. J. Wainscoat\altaffilmark{\IfA} 64 77 % M. F. Waterson,\altaffilmark{\IfA} 65 78 } % this bracket terminates author list 66 79 80 \altaffiltext{\IfA}{Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu HI 96822} 81 \altaffiltext{\LBL}{Lawrence Berkeley National Laboratory, One Cyclotron Road, Berkeley, CA 94720, USA} 82 \altaffiltext{\Hubble}{Hubble Fellow} 83 \altaffiltext{\ITC}{Institute for Theory and Computation, Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, MS-51, Cambridge, MA 02138 USA} 84 \altaffiltext{\Harvard}{Department of Physics, Harvard University, Cambridge, MA 02138 USA} 85 \altaffiltext{\MPIA}{Max Planck Institute for Astronomy, K\"onigstuhl 17, D-69117 Heidelberg, Germany} 86 \altaffiltext{\ARI}{Astronomisches Rechen-Institut, Zentrum f\"ur Astronomie der Universit\"at Heidelberg, M\"ochhofstrasse 12-14, D-69120 Heidelberg, Germany} 87 \altaffiltext{\Princeton}{Department of Astrophysical Sciences, Princeton University, Princeton, NJ 08544, USA} 88 \altaffiltext{\DUR}{Department of Physics, Durham University, South Road, Durham DH1 3LE, UK} 89 \altaffiltext{\CfA}{Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138} 90 67 91 % The ordering here should be sequential, matching the sequence in the list of authors: 68 \altaffiltext{\IfA}{Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu HI 96822}69 % \altaffiltext{\CfA}{Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138}70 % \altaffiltext{\Princeton}{Department of Astrophysical Sciences, Princeton University, Princeton, NJ 08544, USA}71 92 % \altaffiltext{\USNO}{US Naval Observatory, Flagstaff Station, Flagstaff, AZ 86001, USA} 72 93 % \altaffiltext{\JHU}{Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA} 73 % \altaffiltext{\MPIA}{Max Planck Institute for Astronomy, K\"onigstuhl 17, D-69117 Heidelberg, Germany} 94 95 % \altaffiltext{\Strassborg}{ 96 74 97 \begin{abstract} 75 98 76 Lorem ipsum dolor sit amet, consectetur adipiscing elit. Vestibulum 77 bibendum nisi id tristique posuere. Duis eu mollis nulla. Maecenas est 78 turpis, mattis tempor urna vitae, placerat rhoncus sem. Lorem ipsum 79 dolor sit amet, consectetur adipiscing elit. Sed quis velit 80 nisl. Aliquam erat volutpat. Cras lacinia, nisl tristique auctor 81 molestie, dolor nulla rhoncus purus, ac accumsan nunc nunc ac 82 nibh. Maecenas vitae mollis mauris. Ut sollicitudin pulvinar purus, 83 eget luctus lorem tincidunt vitae. Vestibulum eu mattis neque. Nulla 84 in tortor id urna dapibus gravida a vel leo. 99 The Pan-STARRS\,1 $3\pi$ survey has produced photometry and astrometry 100 covering the \approx 30,000 square degrees $\delta > -30$\degrees. 101 This article describes the photometric and astrometric calibration of this survey. 85 102 86 103 \end{abstract} 87 104 88 105 % insert additional keywords as appropriate: 89 %\keywords{Surveys:\PSONE }106 \keywords{Surveys:\PSONE } 90 107 91 108 \section{Introduction}\label{sec:intro} 109 110 This is the fifth in a series of seven papers describing the 111 Pan-STARRS1 Surveys, the data reduction techiques and the resulting 112 data products. This paper (Paper V) describes the final calibration 113 process, and the resulting photometric and astrometric quality. 114 115 %Chambers et al. 2017 (Paper I) 116 %The Pan-STARRS\,1 Surveys 117 \citet[][Paper I]{chambers2017} 118 provides an overview of the Pan-STARRS System, the design and 119 execution of the Surveys, the resulting image and catalog data 120 products, a discussion of the overall data quality and basic 121 characteristics, and a brief summary of important results. 122 123 %Magnier et al. 2017 (Paper II) 124 %Pan-STARRS Data Processing Stages 125 \citet[][Paper II]{magnier2017c} 126 describes how the various data processing stages are organised and implemented 127 in the Imaging Processing Pipeline (IPP), including details of the 128 the processing database which is a critical element in the IPP infrastructure . 129 130 %Waters et al. 2017 (Paper III) 131 %Pan-STARRS Pixel Processing : Detrending, Warping, Stacking 132 \citet[][Paper III]{waters2017} 133 describes the details of the pixel processing algorithms, including detrending, warping, and adding (to create stacked images) and subtracting (to create difference images) and resulting image products and their properties. 134 135 136 %Magnier et al. 2017 (Paper IV) 137 %Pan-STARRS Pixel Analysis : Source Detection 138 \citet[][Paper IV]{magnier2017a} 139 describes the details of the source detection and photometry, including point-spread-function and extended source fitting models, and the techniques for ``forced" photometry measurements. 140 141 %Magnier et al. 2017 (Paper V) 142 %Pan-STARRS Photometric and Astrometric Calibration 143 %\citet[][Paper V]{magnier2017b} 144 %describes the final calibration process, and the resulting photometric and astrometric quality. 145 146 147 %Flewelling et al. 2017 (Paper VI) 148 %Pan-STARRS 1 Database and Data Products 149 \citet[][Paper VI]{flewelling2017} 150 describes the details of the resulting catalog data and its organization in the Pan-STARRS database. 151 % 152 % 153 \citet[][Paper VII]{huber2017} 154 %Huber et al. 2017 (Paper VII) 155 describes the Medium Deep Survey in detail, including the unique issues and data products specific to that survey. The Medium Deep Survey is not part of Data Release 1. (DR1) 156 157 % 158 The Pan-STARRS1 filters and photometric system have already been 159 described in detail in \cite{2012ApJ...750...99T}. 160 161 {\color{red} {\em Note: These papers are being placed on arXiv.org to 162 provide crucial support information at the time of the public 163 release of Data Release 1 (DR1). We expect the arXiv versions to 164 be updated prior to submission to the Astrophysical Journal in 165 January 2017. Feedback and suggestions for additional information 166 from early users of the data products are welcome during the 167 submission and refereeing process.}} 92 168 93 169 \section{Pan-STARRS\,1} … … 103 179 The wide-field \PSONE\ telescope consists of a 1.8~meter diameter 104 180 $f$/4.4 primary mirror with an 0.9~m secondary, producing a 3.3 degree 105 field of view \citep{ PS1.optics}. The optical design yields low181 field of view \citep{2004SPIE.5489..667H}. The optical design yields low 106 182 distortion and minimal vignetting even at the edges of the illuminated 107 183 region. The optics, in combination with the natural seeing, result in 108 generally good image quality: 75\% of the images have full-width 109 half-max values less than \note{(1.X, 1.X, 1.X, 1.X, 1.X), update} 110 arcseconds for (\grizy), with a floor of $\sim 0.7$ \note{update} 111 arcseconds. The \PSONE\ camera \citep{PS1.GPCA} is a mosaic of 60 112 edge-abutted $4800\times4800$ pixel back-illuminated \note{name} CCDs 113 manufactured by Lincoln Laboratory. The CCDs have 10~$\mu$m pixels 114 subtending 0.258~arcsec and are \note{70um} thick. The detectors are 115 read out using a StarGrasp CCD controller, with a readout time of 7 116 seconds for a full unbinned image \citep{PS1.GPCB}. The active, 117 usable pixels cover $\sim 80$\% of the FOV. 184 generally good image quality: the median image quality for the 3$\pi$ 185 survey is FWHM = (1.31, 1.19, 1.11, 1.07, 1.02) arcseconds for 186 (\grizy), with a floor of $\sim0.7$ arcseconds. The \PSONE\ camera 187 \citep{PS1.GPCA} is a mosaic of 60 edge-abutted $4800\times4800$ pixel 188 back-illuminated CCID58 Orthogonal Transfer Arrays manufactured by 189 Lincoln Laboratory \citep{2006amos.confE..47T,2008SPIE.7021E..05T}. 190 The CCDs have 10~$\mu$m pixels subtending 0.258~arcsec and are 191 70$\mu$m thick. The detectors are read out using a StarGrasp CCD 192 controller, with a readout time of 7 seconds for a full unbinned image 193 \citep{2008SPIE.7014E..0DO}. The active, usable pixels cover $\sim 194 80$\% of the FOV. 118 195 119 196 Nightly observations are conducted remotely from the Advanced … … 127 204 128 205 Images obtained by \PSONE\ are automatically processed in real time by 129 the \PSONE\ Image Processing Pipeline \citep[IPP,][]{ PS1.IPP}.206 the \PSONE\ Image Processing Pipeline \citep[IPP,][]{magnier2017a}. 130 207 Real-time analysis goals are aimed at feeding the discovery pipelines 131 208 of the asteroid search and supernova search teams. The data obtained … … 196 273 \section{Astrometric Models} 197 274 275 % \note{include projection math?} 276 % \note{reference discussion somewhere on cell vs chip} 277 198 278 Three somewhat distinct astrometric models are employed within the IPP 199 279 at different stages. The simplest model is defined independently for 200 280 each chip: a simple TAN projection (Calabretta \& Griesen REF) is used 201 281 to relate sky coordinates to a cartesian tangent-plane coordinate 202 system. \note{include projection math?}A pair of low-order282 system. A pair of low-order 203 283 polynomials are used to relate the chip pixel coordinates to this 204 284 tangent-plane coordinate system. The transforming polynomials are of … … 209 289 \end{eqnarray} 210 290 where $P,Q$ are the tangent plane coordinates, $X_{\rm chip}, Y_{\rm 211 chip}$ are the coordinates on the 60 GPC1 chips (\note{see 212 discussion somewhere on cell vs chip}), and $C^P_{i,j}, C^Q_{i,j}$ 291 chip}$ are the coordinates on the 60 GPC1 chips, and $C^P_{i,j}, C^Q_{i,j}$ 213 292 are the polynomial coefficients for each order. In the \code{psastro} 214 293 analysis, $i + j <= N_{\rm order}$ where the order of the fit, $N_{\rm 215 294 order}$, may be 1 to 3, under the restriction that sufficient stars 216 are needed to constraint the order \note{describe a bit better: this 217 is automatically selected based on the number of stars}. 295 are needed to constrain the order. 296 297 % \note{describe a bit better: this is automatically selected based on the number of stars} 218 298 219 299 A second form of astrometry model which yields somewhat higher … … 234 314 code restricts the exponents with the rule $i + j <= N_{\rm order}$ 235 315 where the order of the fit, $N_{\rm order}$, may be 1 to 3, under the 236 restriction that sufficient stars are needed to constraint the order 237 \note{describe a bit better: this is automatically selected based on 238 the number of stars}. 316 restriction that sufficient stars are needed to constrain the order 239 317 For each chip, a second set of polynomials describes the 240 318 transformation from the chip coordinate systems to the focal … … 270 348 Q & = & \sum_{i,j} C^Q_{i,j} (X_{\rm chip} - X_0)^i (Y_{\rm chip} - Y_0)^j 271 349 \end{eqnarray} 272 \note{need to complete this discussion of the WCS keywords, both 273 standard and non-standard, used to represent these polynomial 274 transformations} 275 276 \begin{verbatim} 277 Here is a table of the keywords and the related terms from Eqns above: 278 CTYPE1,2 : RA---WRP, DEC--WRP & RA---DIS, DEC--DIS 279 CRVAL1,2 : C^{L,M}_{0,0} 280 CRPIX1,2 : X_0, Y_0 281 PC001001 : C^{L}_{1,0} 282 PC001002 : C^{L}_{0,1} 283 PC002001 : C^{M}_{1,0} 284 PC002002 : C^{M}_{0,1} 285 PCA1XiYj : C^{L}_{i,j} 286 PCA2XiYj : C^{M}_{i,j} 287 \end{verbatim} 350 351 %% \note{need to complete this discussion of the WCS keywords, both 352 %% standard and non-standard, used to represent these polynomial 353 %% transformations} 354 355 %% \begin{verbatim} 356 %% Here is a list of the keywords 357 %% and the related terms from Eqns above: 358 %% CTYPE1,2 : RA---WRP, DEC--WRP 359 %% CTYPE1,2 : RA---DIS, DEC--DIS 360 %% CRVAL1,2 : C^{L,M}_{0,0} 361 %% CRPIX1,2 : X_0, Y_0 362 %% PC001001 : C^{L}_{1,0} 363 %% PC001002 : C^{L}_{0,1} 364 %% PC002001 : C^{M}_{1,0} 365 %% PC002002 : C^{M}_{0,1} 366 %% PCA1XiYj : C^{L}_{i,j} 367 %% PCA2XiYj : C^{M}_{i,j} 368 %% \end{verbatim} 288 369 289 370 \section{Real-time Calibration} … … 318 399 reference catalog generated from internal re-calibration of the PV0 319 400 analysis of PS1 photometry and astrometry was used for the reference 320 catalog. \note{discuss history of the different refcats?} 401 catalog. 402 403 % \note{discuss history of the different refcats?} 321 404 322 405 Coordinates and calibrated magnitudes of stars from the reference … … 326 409 position angle reported by the header. Reference stars are selected 327 410 from the full field of view of the GPC1 camera, padded by an 328 additional \note{25\%}to ensure a match can be determined even in the411 additional 25\% to ensure a match can be determined even in the 329 412 presence of substantial errors in the boresite coordinates. It is 330 413 important to choose an appropriate set of reference stars: if too few … … 366 449 \end{eqnarray} 367 450 are generated. The collection of $\Delta X, \Delta Y$ values are 368 collected in a 2D histogram with sampling of \note{XXX}pixels and the451 collected in a 2D histogram with sampling of 50 pixels and the 369 452 peak pixel is identified. If the astrometry guess were perfect, this 370 453 peak pixel would be expected to lie at (0,0) and contain all of the … … 391 474 astrometry guess for the chip. 392 475 393 \note{option to downweight based on photometric inconsistency : not 394 used in PS1 analysis} 476 %% \note{option to downweight based on photometric inconsistency : not used in PS1 analysis} 395 477 396 478 \subsection{Chip Polynomial Fits} … … 435 517 desired for the distortion fit. The coefficients of the gradient fit 436 518 are then used to determine the coefficients for the polynomials 437 representing the distortion. \note{write out the math of the gradients} 519 representing the distortion. 520 521 %% \note{write out the math of the gradients} 438 522 439 523 Once the common distortion coming from the optics and atmosphere have 440 524 been modeled, \code{psastro} determines polynomial transformations 441 525 from the 60 chips to the focal plane coordinate system. In this 442 stage, \note{NN} iterations of the chip fits are performed. Before 443 each iteration, the reference stars and detected objects are matched 444 using the current best set of transformations. These fits start with 445 low order (1) and large matching radius (\note{XX}). As the 446 iterations proceed, the radius is reduced and the order is allowed to 447 increaes, up to 3rd order for the final iterations. \note{quality of 448 the fits as a result of this stage}. 526 stage, 5 iterations of the chip fits are performed. Before each 527 iteration, the reference stars and detected objects are matched using 528 the current best set of transformations. These fits start with low 529 order (1) and large matching radius. As the iterations proceed, the 530 radius is reduced and the order is allowed to increaes, up to 3rd 531 order for the final iterations. 532 533 %% \note{quality of the fits as a result of this stage}. 449 534 450 535 \subsection{Real-time Photometric Calibration} 536 537 %% \note{define / describe the robust median} 451 538 452 539 After the astrometric calibration has finished, the photometric 453 540 calibration is performed by \code{psastro}. When the reference stars 454 541 are loaded, the apparent magnitude in the filter of interest is also 455 loaded. 542 loaded. Stars for which the reference magnitude is brighter than 456 543 (\grizy) = (19, 19, 18.5, 18.5, 17.5) are used to determine the zero 457 544 points by comparison with the instrumental magnitudes. For the PV3 458 analysis, the robust median \note{defined where?} is used to measure459 the zero point. For early versions of the analysis, when the reference 460 catalog used synthetic magnitudes, it was necessary to search for the461 blue edge of the distribution: the synthetic magnitude poorly 462 predicted the magnitudes of stars in the presence of significant 463 extinction or for the very red stars, making the blue edge somewhat 464 more reliable. Note that we do not include an airmass correction in 465 this zero point analysis: the airmass correction is folded into the 466 observed zero point. The zero point may be measured separately for 467 each chip or as a single value for the entire exposure; the latter468 option was used forthe PV3 analysis.545 analysis, an outlier-rejecting median is used to measure the zero 546 point. For early versions of the analysis, when the reference catalog 547 used synthetic magnitudes, it was necessary to search for the blue 548 edge of the distribution: the synthetic magnitude poorly predicted the 549 magnitudes of stars in the presence of significant extinction or for 550 the very red stars, making the blue edge somewhat more reliable. Note 551 that we do not include an airmass correction in this zero point 552 analysis: the airmass correction is folded into the observed zero 553 point. The zero point may be measured separately for each chip or as 554 a single value for the entire exposure; the latter option was used for 555 the PV3 analysis. 469 556 470 557 \subsection{Real-time outputs} … … 483 570 chip-level keywords (e.g., \code{DATE-OBS}). The astrometric 484 571 transformation information for each chip is saved in the corresponding 485 header using standard (and some non-standard) WCS keywords. 486 \note{combine this discussion with the above?}. For the two-level 487 astrometric model, the PHU header carries the astrometric 572 header using standard (and some non-standard) WCS keywords. For the 573 two-level astrometric model, the PHU header carries the astrometric 488 574 transformation related to the projection and the camera-wide 489 575 distortions. Photometric calibrations are written as a set of … … 507 593 \subsection{Ubercal Analysis} 508 594 509 \note{clean up and re-word the pieces below}595 % \note{clean up and re-word the pieces below} 510 596 511 597 The photometric calibration of the DVO database starts with the 512 ``ubercal'' analysis technique as described by \cite{ PS1.ubercal}.598 ``ubercal'' analysis technique as described by \cite{2012ApJ...756..158S}. 513 599 This analysis is performed by the group at Harvard, loading data from 514 600 the \code{smf} files into their instance of the Large Scale Database … … 517 603 518 604 Photometric nights are selected and all other exposures are ignored. 519 Each night \note{shorter time?} is allowed to have a single fitted520 zero point and a single fitted value for the airmass extinction 521 coefficient per filter. The zero points and extinction terms are 522 determined as a least squares minimization process using the repeated 523 measurements of the same stars from different nights to tie nights524 together. Flat-field corrections are also determined as part ofthe525 minimization process. In the original (PV1) ubercal analysis, 526 \cite{PS1.ubercal} determined flat-field corrections for $2\times 2$ 527 sub-regions of each chip in the camera and four distinct time periods528 ( ``seasons''). Later analysis (PV2) used an $8\times8$ grid of529 flat-field corrections to goodeffect.605 Each night is allowed to have a single fitted zero point and a single 606 fitted value for the airmass extinction coefficient per filter. The 607 zero points and extinction terms are determined as a least squares 608 minimization process using the repeated measurements of the same stars 609 from different nights to tie nights together. Flat-field corrections 610 are also determined as part of the minimization process. In the 611 original (PV1) ubercal analysis, \cite{2012ApJ...756..158S} determined 612 flat-field corrections for $2\times 2$ sub-regions of each chip in the 613 camera and four distinct time periods (``seasons''). Later analysis 614 (PV2) used an $8\times8$ grid of flat-field corrections to good 615 effect. 530 616 531 617 The ubercal analysis was re-run for PV3 by the Harvard group. For the … … 536 622 was also included for PV3. In retrospect, as we show below, the data 537 623 from the latter part of the survey would probably benefit from 538 additional flat-field seasons. \note{something for PV4}. 624 additional flat-field seasons. 625 626 %% \note{something for PV4}. 539 627 540 628 By excluding non-photometric data and only fitting 2 parameters for … … 545 633 every night, helping to tie down overall variations of the system 546 634 throughput and acting as internal standard star fields. The resulting 547 photometric system is shown by \cite{ PS1.ubercal} to have reliability635 photometric system is shown by \cite{2012ApJ...756..158S} to have reliability 548 636 across the survey region at the level of (8.0, 7.0, 9.0, 10.7, 12.4) 549 637 millimags in (\grizy). As we discuss below, this conclusion is 550 reinforced by our external comparison. \note{do I have a measurement 551 of the bright end stability in PV3? basically, what is the scatter 552 per star as a function of position in the camera and magnitude?} 638 reinforced by our external comparison. 639 640 %% \note{do I have a measurement 641 %% of the bright end stability in PV3? basically, what is the scatter 642 %% per star as a function of position in the camera and magnitude?} 553 643 554 644 The overall zero point for each filter is not naturally determined by 555 645 the Ubercal analysis; an external constraint on the overall 556 photometric system is required for each filter. \cite{PS1.ubercal} 557 used photometry of the MD09 Medium Deep field to match the photometry 558 measured by \cite{JTphoto} on the reference photometric night of MJD 559 55744 (UT 02 July 2011). \note{Scolnic et al REF} have re-examined 560 the photometry of Calspec standards as observed by PS1. They reject 2 561 of the \note{XX} stars used by \cite{JTphoto} and add photometry of 562 \note{XX} additional stars. The calspec spectrophotometry values have 563 also been re-examined by XX; using these new measurements, Scolnic et 564 al determine new zero points for the PS1 system, which we have applied 565 (see below). 646 photometric system is required for each filter. 647 \cite{2012ApJ...756..158S} used photometry of the MD09 Medium Deep 648 field to match the photometry measured by \cite{2012ApJ...750...99T} 649 on the reference photometric night of MJD 55744 (UT 02 July 2011). 650 \cite{2015ApJ...815..117S} have re-examined the photometry of Calspec 651 standards as observed by PS1. They reject 2 of the 5 stars used by 652 \cite{2012ApJ...750...99T} and add photometry of 2 additional stars. 653 654 %% \note{The calspec spectrophotometry values have also been re-examined 655 %% by REF; using these new measurements, \cite{2015ApJ...815..117S} 656 %% determine new zero points for the PS1 system, which we have applied 657 %% (see below).} 566 658 567 659 \subsection{Applying the Ubercal Zero Points : Setphot} … … 585 677 each filter representing respectively the nominal zero point and the 586 678 slope of the trend with respect to the airmass ($\zeta$) for each 587 filter. \note{the image zero point does not incorporate the airmass,588 only the measurement zero point}. These static values are listed in 589 Table~\ref{tab:zpts}. When \code{setphot} was run, these static zero 590 points have been adjusted by the calspec offsets listed in 591 Table~\ref{tab:zpts} based on the analysis of CALSPEC standards by 592 Scolnic et al REF. These offsets bring the photometric system defined 593 by the ubercal analysis into alignment with the Scolnic analysis of 594 the PS1 observations of XXX calspec standard stars. The value 595 $M_{cal}$ is the offset needed by each exposure to match the ubercal 596 value, or to bring the non-ubercal exposures into agreement with the 597 rest of the exposures, as discussed below. The flat-field information 598 i s encoded in a table of flat-field offsets as a function of time,599 fi lter, and camera position. Each image which is part of the ubercal600 subset is marked with a bit in the field \code{Image.flags}: 601 \code{ID_IMAGE_PHOTOM_UBERCAL = 0x00000200} 679 filter. These static values are listed in Table~\ref{tab:zpts}. When 680 \code{setphot} was run, these static zero points have been adjusted by 681 the calspec offsets listed in Table~\ref{tab:zpts} based on the 682 analysis of CALSPEC standards by Scolnic et al REF. These offsets 683 bring the photometric system defined by the ubercal analysis into 684 alignment with the Scolnic analysis of the PS1 observations of XXX 685 calspec standard stars. The value $M_{cal}$ is the offset needed by 686 each exposure to match the ubercal value, or to bring the non-ubercal 687 exposures into agreement with the rest of the exposures, as discussed 688 below. The flat-field information is encoded in a table of flat-field 689 offsets as a function of time, filter, and camera position. Each 690 image which is part of the ubercal subset is marked with a bit in the 691 field \code{Image.flags}: \code{ID_IMAGE_PHOTOM_UBERCAL = 0x00000200} 692 693 %% \note{give airmass formula for completeness?}. 602 694 603 695 When \code{setphot} applies the ubercal information to the image … … 611 703 with the airmass for the measurement, calculated using the altitude of 612 704 the individual detection as determined from the Right Ascension, 613 Declination, the observatory latitude, and the sidereal time. 614 \note{give formula for completeness?}. For a camera with the field of 615 view of the PS1 GPC1, the airmass may vary significantly within the 616 field of view, especially at low elevations. In the worst cases, at 617 the celestial pole, the airmass range within a single exposure is XXX 618 - XXX. The complete calibrated (`relative') magnitude is determined 619 from the stored database values as: 705 Declination, the observatory latitude, and the sidereal time. For a 706 camera with the field of view of the PS1 GPC1, the airmass may vary 707 significantly within the field of view, especially at low elevations. 708 In the worst cases, at the celestial pole, the airmass range within a 709 single exposure is XXX - XXX. The complete calibrated (`relative') 710 magnitude is determined from the stored database values as: 620 711 \[ 621 712 M_{\rm rel} = M_{\rm inst} - 25.0 + zp_{\rm ref} + M_{\rm cal} + M_{\rm flat} + K_\lambda (sec \zeta - 1). … … 643 734 \subsection{Relphot Analysis} 644 735 736 %% \note{how many exposures are not in ubercal?} 737 645 738 Relative photometry is used to determine the zero points of the 646 exposures which were not included in the ubercal analysis \note{how 647 many?}. The relative photometry analysis has been desribed in the 739 exposures which were not included in the ubercal analysis. The relative photometry analysis has been desribed in the 648 740 past in Magnier et al 2013 REF. We review that analysis here, along 649 741 with specific updates for PV3. … … 660 752 \[ M_{ave} = \frac{\sum_i M_{rel,i} w_i}{\sum_i w_i} \] 661 753 We find that the color difference of the different chips can be 662 ignored \note{level of this effect?}, and set the value of $A$ to 0.0.754 ignored, and set the value of $A$ to 0.0. 663 755 Note that we only use a single mean airmass extinction term for all 664 756 exposures -- the difference between the mean and the specific value 665 757 for a given night is taken up as an additional element of the 666 758 atmospheric attenuation. 759 760 %% \note{color-color terms between chips?} 667 761 668 762 We write a global $\chi^2$ equation which we attempt to minimize by … … 681 775 Only brighter, high quality measurements are used in the relative 682 776 photometry analysis of the exposure zero points. We use only the 683 brighter objects \note{mag limit}, limiting the density to a maximum 684 of \note{actual max density?} 2500 or 3000 objects per square degree 685 (lower in areas where we have more observations). When limiting the 686 density, we prefer objects which are brighter (but not saturated), and 687 those with the most measurements (to ensure better coverage over the 688 available images). 777 brighter objects, limiting the density to a maximum of 4000 objects 778 per square degree (lower in areas where we have more observations). 779 When limiting the density, we prefer objects which are brighter (but 780 not saturated), and those with the most measurements (to ensure better 781 coverage over the available images). 689 782 690 783 There are a few classes of outliers which we need to be careful to … … 694 787 We attempt to exclude these poor measurements in advance by rejecting 695 788 measurements which the photometric analysis has flagged the result as 696 suspcious. \note{bad and poor psphot bits?} We reject detections 697 which are excessively masked ({\tt PSF\_QF} $<$ 0.85, see Magnier et 698 al PSPHOT REF); these include detections which are too close to other 699 bright objects, diffraction spikes, ghost images, or the detector 700 edges. However, these rejections do not catch all cases of bad 701 measurements. 789 suspcious. We reject detections which are excessively masked; these include 790 detections which are too close to other bright objects, diffraction 791 spikes, ghost images, or the detector edges. However, these 792 rejections do not catch all cases of bad measurements. 793 794 %% \citep[\code{PSF_QF} $< 0.85$, see][]{magnier2017b}; 795 %% \note{refer to the PSPHOT bad and poor psphot bits?} 702 796 703 797 After the initial iterations, we also perform outlier rejections based … … 713 807 with reduced $\chi^2$ values more than 20.0, or more than 2$\times$ 714 808 the median, whichever is larger. We also exclude stars with standard 715 deviation (of the measurements used for the mean) greater than 716 \note{is this true?} 0.005 mags or 2$\times$ the median standard 717 deviation, whichever is greater. 809 deviation (of the measurements used for the mean) greater than 0.005 810 mags or 2$\times$ the median standard deviation, whichever is greater. 811 812 %% \note{is this true?} 718 813 719 814 Similarly for images, we exclude those with more than 2 magnitudes of … … 734 829 calculation of the formal error on the mean magnitudes propagates this 735 830 additional weight, so that the errors on the Ubercal observations 736 dominates where they are present. \note{do we drop this when 737 calculating the final mean mags?} 831 dominates where they are present. 832 833 % \note{do we drop this when calculating the final mean mags?} 738 834 % \note{do I need to present the math?} 739 835 \[ \mu = \frac{\sum m_i w_i \sigma_i^{-2}}{\sum w_i \sigma_i^{-2}} \] … … 802 898 analysis. 803 899 804 \note{need to discuss the process of setting the final mean magnitudes}900 %% \note{need to discuss the process of setting the final mean magnitudes} 805 901 806 902 For PV3, the relphot analysis was performed two times. The first … … 812 908 data in DVO after the initial relphot calibration to measure the 813 909 flat-field residual with much finer resolution: 124 x 124 flat-field 814 values for each GPC1 chip (40x40 pixels per point). \note{show the 815 flat-field residual images, discuss the features?}. We then used 910 values for each GPC1 chip (40x40 pixels per point). We then used 816 911 \code{setphot} to apply this new flat-field correction, as well as the 817 912 ubercal flat-field corrections, to the data in the database. At this 818 913 point, we re-ran the entire relphot analysis to determine zero points 819 914 and to set the average magnitudes. 915 916 %% \note{show the flat-field residual images, discuss the features?}. 820 917 821 918 For stacks and warps, the image calibrations were determined after the … … 831 928 appropriate for a given warp. This latter effect is one of several 832 929 which degrade the warp photometry compared to the chip photometry at 833 the bright end. \note{recommendation} 930 the bright end. 931 932 %% \note{recommendation} 834 933 835 934 \subsection{Calculation of Object Photometry} … … 859 958 the master DVO database. 860 959 861 \note{need to describe the assignment of flags, etc, for the external 862 data sources}. 960 %% \note{need to describe the assignment of flags, etc, for the external data sources}. 863 961 864 962 \section{Astrometry Analysis} … … 915 1013 contaminated by the effect. 916 1014 1015 % \note{was there is significant difference using a surface brightness version?} 1016 917 1017 We measured the Koppenh\"offer Effect by accumulating the residual 918 astrometry statistics for \note{how many} stars. For each chip, we1018 astrometry statistics for stars in the database. For each chip, we 919 1019 measured the mean X and Y displacements of the astrometric residuals 920 1020 as function of the instrumental magnitude of the star divided by the 921 FWHM$^2$. \note{was there is significant difference using a surface 922 brightness version?} We measured the trend for all chips in a 1021 FWHM$^2$. We measured the trend for all chips in a 923 1022 number of different time ranges and found the effect to be quite 924 1023 stable, in the period where it was present. The effect only appeared … … 964 1063 DCR trend for the 5 filters \grizy, as well as the measured 965 1064 displacement in the direction perpendicular to the parallactic angle. 966 We represent the trend with a spline fitted to this dataset. The DCR967 trend has an amplitude of \note{XXX - XXX} in the five filters. 968 969 \note{write down the DCR formalae for reference}.1065 We represent the trend with a spline fitted to this dataset. 1066 1067 %% The DCR trend has an amplitude of \note{XXX - XXX} in the five filters. 1068 %% \note{write down the DCR formalae for reference}. 970 1069 971 1070 \subsubsection{Astrometric Flat-field} … … 1017 1116 similar to the ``tree rings'' reported by the DES team and others 1018 1117 (G. Berstein REF \& REFS). We explore these tree rings in detail in 1019 \note{SECTION or REF?}. 1118 1119 % \note{SECTION or REF?}. 1020 1120 1021 1121 After the initial analysis to measure the KE corrections, DCR … … 1091 1191 performed SED fitting for 800M stars in the 3$\pi$ region using PV2 1092 1192 data. The goal of this work was to determine the 3D structure of the 1093 dust in the galaxy. By fitting model SEDs to \note{all?} stars1094 meeting a basic data quality cut \note{(describe)}, they determined 1095 the best spectral type, and thus $T_{\rm eff}$, absolute $r$-band1096 magnitude, distance modulus, and extinction $A_V$ (the desired output1097 and used to determine the dust extinction as a function of distance1098 th roughout the galaxy). We use the distance modulus determined in1099 this analysis to predict the proper motions. 1193 dust in the galaxy. By fitting model SEDs to stars meeting a basic 1194 data quality cut, they determined the best spectral type, and thus 1195 $T_{\rm eff}$, absolute $r$-band magnitude, distance modulus, and 1196 extinction $A_V$ (the desired output and used to determine the dust 1197 extinction as a function of distance throughout the galaxy). We use 1198 the distance modulus determined in this analysis to predict the proper 1199 motions. 1100 1200 1101 1201 To convert the distances to proper motions, we use the Galactic … … 1112 1212 \end{eqnarray} 1113 1213 where $d$ is the distance and $l,b$ are the Galactic coordintes of the 1114 star. \note{some reference?} Note that the proper motion induced by 1214 star. Note that the proper motion induced by 1215 %% \note{some reference for this?} 1115 1216 the Galactic rotation is independent of distance while the reflex 1116 1217 motion induced by the solar motion decreases with increasing … … 1126 1227 value of 500pc. 1127 1228 1128 \note{plots to show how well this worked for PV3 pre Gaia}1229 %% \note{plots to show how well this worked for PV3 pre Gaia} 1129 1230 1130 1231 \subsection{Gaia Constraint} … … 1136 1237 motion and parallax solutions are in general saturated in the PS1 1137 1238 observations. Thus, we are limited to using the Gaia mean positions 1138 reported for the fainter stars. We extracted all Gaia sources 1139 \note{not marked as a duplicate} from \note{where?} and generated a 1140 DVO database from this dataset. We then merged the Gaia DVO into the 1141 PV3master DVO database. We re-ran the complete relative astrometry1239 reported for the fainter stars. We extracted all Gaia sources not 1240 marked as a duplicate from the Gaia archive and generated a DVO 1241 database from this dataset. We then merged the Gaia DVO into the PV3 1242 master DVO database. We re-ran the complete relative astrometry 1142 1243 analysis using Gaia as an additional measurement. We applied the 1143 1244 analysis described above, applying the estimated distances to … … 1154 1255 even at a lower weight, helps to tile over those gaps. 1155 1256 1156 \note{Figures showing the Gaia residuals}1257 %% \note{Figures showing the Gaia residuals} 1157 1258 1158 1259 \subsection{Calculation of Object Astrometry} … … 1165 1266 1166 1267 \section{Conclusion} 1268 1269 \acknowledgments 1270 1271 The Pan-STARRS1 Surveys (PS1) have been made possible through 1272 contributions of the Institute for Astronomy, the University of 1273 Hawaii, the Pan-STARRS Project Office, the Max-Planck Society and its 1274 participating institutes, the Max Planck Institute for Astronomy, 1275 Heidelberg and the Max Planck Institute for Extraterrestrial Physics, 1276 Garching, The Johns Hopkins University, Durham University, the 1277 University of Edinburgh, Queen's University Belfast, the 1278 Harvard-Smithsonian Center for Astrophysics, the Las Cumbres 1279 Observatory Global Telescope Network Incorporated, the National 1280 Central University of Taiwan, the Space Telescope Science Institute, 1281 the National Aeronautics and Space Administration under Grant 1282 No. NNX08AR22G issued through the Planetary Science Division of the 1283 NASA Science Mission Directorate, the National Science Foundation 1284 under Grant No. AST-1238877, the University of Maryland, and Eotvos 1285 Lorand University (ELTE) and the Los Alamos National Laboratory. 1286 1287 \bibliographystyle{apj} 1288 %\bibliography{lib}{} 1289 \input{calibration.bbl} 1290 1291 \end{document} 1167 1292 1168 1293 \begin{verbatim} … … 1184 1309 \end{verbatim} 1185 1310 1186 \end{document} -
trunk/doc/release.2015/ps1.datasystem/Makefile
r39848 r39868 1 1 # $Id: Makefile,v 1.16 2006-01-16 01:11:40 eugene Exp $ 2 3 DO_PDFLATEX = 1 4 DO_BIBTEX = 1 2 5 3 6 help: 4 7 @echo "USAGE: make (target)" 5 @echo " targets: all datasystem"8 @echo " targets: all tgz pdf" 6 9 7 all: datasystem.pdf 8 datasystem: datasystem.pdf 10 all: pdf tgz 11 tgz: datasystem.tgz 12 pdf: datasystem.pdf 9 13 10 DATASYSTEM = datasystem.tex 14 FILES = \ 15 ../inputs/astro.sty \ 16 ../inputs/code.sty \ 17 ../inputs/apj.bst \ 18 ../inputs/lib.bib \ 19 datasystem.tex \ 20 datasystem.bbl 11 21 12 # pics/Metadata.ps 13 # pics/earthrot.ps 14 15 datasystem.pdf: $(DATASYSTEM) 16 17 datasystem.ps: $(DATASYSTEM) 22 datasystem.pdf: $(FILES) 23 datasystem.tgz: $(FILES) 18 24 19 25 include ../Makefile.Common -
trunk/doc/release.2015/ps1.datasystem/datasystem.tex
r39865 r39868 1 %\documentclass[iop,floatfix]{emulateapj}1 \documentclass[iop,floatfix]{emulateapj} 2 2 % \documentclass[iop,floatfix,onecolumn]{emulateapj} 3 \documentclass[12pt,preprint]{aastex}3 % \documentclass[12pt,preprint]{aastex} 4 4 % \pdfoutput=1 5 5 … … 29 29 % list and (2) re-order the list at the bottom (and comment-out as needed) 30 30 \def\IfA{1} 31 \def\CfA{2} 32 \def\MPIA{3} 31 \def\LSST{2} 33 32 \def\Princeton{3} 34 \def\ USNO{4}35 \def\ JHU{1}33 \def\DUR{4} 34 \def\CfA{5} 36 35 37 36 % This example has a first author from UH: 38 37 \author{ 39 38 Eugene A. Magnier,\altaffilmark{\IfA} 40 IPP Team, 39 K.~C. Chambers,\altaffilmark{\IfA} 40 H.~A. Flewelling,\altaffilmark{\IfA} 41 J.~C. Hoblitt,\altaffilmark{\LSST} 42 M. E. Huber,\altaffilmark{\IfA} 43 % R. H. Lupton,\altaffilmark{\Princeton} 44 P.~A. Price,\altaffilmark{\Princeton} 45 W.~E. Sweeney,\altaffilmark{\IfA} 46 C. Z. Waters,\altaffilmark{\IfA} 41 47 %PS Builder List 48 L. Denneau,\altaffilmark{\IfA} 49 P. Draper,\altaffilmark{\DUR} 50 K. W. Hodapp,\altaffilmark{\IfA} 51 R. Jedicke,\altaffilmark{\IfA} 52 R.-P. Kudritzki,\altaffilmark{\IfA} 53 N. Metcalfe,\altaffilmark{\DUR} 54 C.~W. Stubbs,\altaffilmark{\CfA} 42 55 % W.~S. Burgett,\altaffilmark{\IfA} 43 % K.~C. Chambers,\altaffilmark{\IfA}44 % L. Denneau,\altaffilmark{\IfA}45 % P. Draper,\altaffilmark{\DUR}46 % H.~A. Flewelling,\altaffilmark{\IfA}47 56 % T. Grav,\altaffilmark{\IfA} 48 57 % J. N. Heasley,\altaffilmark{\IfA} 49 % K. W. Hodapp,\altaffilmark{\IfA}50 % M. E. Huber,\altaffilmark{\IfA}51 % R. Jedicke,\altaffilmark{\IfA}52 58 % N. Kaiser,\altaffilmark{\IfA} 53 % R.-P. Kudritzki,\altaffilmark{\IfA}54 59 % G. A. Luppino,\altaffilmark{\IfA} 55 % R. H. Lupton,\altaffilmark{\Princeton}56 % E. A. Magnier,\altaffilmark{\IfA}57 % N. Metcalfe,\altaffilmark{\DUH}58 60 % D. G. Monet,\altaffilmark{\USNO} 59 61 % J.~S. Morgan,\altaffilmark{\IfA} 60 62 % P. M. Onaka,\altaffilmark{\IfA} 61 % P.~A. Price,\altaffilmark{\Princeton}62 % C.~W. Stubbs,\altaffilmark{\CfA}63 % W.~E. Sweeney,\altaffilmark{\IfA}64 63 % J.~L. Tonry, \altaffilmark{\IfA} 65 % R. J. Wainscoat,\altaffilmark{\IfA} and 66 % C. Z. Waters,\altaffilmark{\IfA} 64 R. J. Wainscoat\altaffilmark{\IfA} 67 65 } % this bracket terminates author list 68 66 69 67 % The ordering here should be sequential, matching the sequence in the list of authors: 70 68 \altaffiltext{\IfA}{Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu HI 96822} 71 % \altaffiltext{\CfA}{Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138} 69 \altaffiltext{\LSST}{LSST Project Management Office, Tucson, AZ, U.S.A} 70 \altaffiltext{\Princeton}{Department of Astrophysical Sciences, Princeton University, Princeton, NJ 08544, USA} 71 \altaffiltext{\DUR}{Department of Physics, Durham University, South Road, Durham DH1 3LE, UK} 72 \altaffiltext{\CfA}{Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138} 72 73 % \altaffiltext{\Princeton}{Department of Astrophysical Sciences, Princeton University, Princeton, NJ 08544, USA} 73 74 % \altaffiltext{\USNO}{US Naval Observatory, Flagstaff Station, Flagstaff, AZ 86001, USA} … … 76 77 \begin{abstract} 77 78 78 Lorem ipsum dolor sit amet, consectetur adipiscing elit. Vestibulum 79 bibendum nisi id tristique posuere. Duis eu mollis nulla. Maecenas est 80 turpis, mattis tempor urna vitae, placerat rhoncus sem. Lorem ipsum 81 dolor sit amet, consectetur adipiscing elit. Sed quis velit 82 nisl. Aliquam erat volutpat. Cras lacinia, nisl tristique auctor 83 molestie, dolor nulla rhoncus purus, ac accumsan nunc nunc ac 84 nibh. Maecenas vitae mollis mauris. Ut sollicitudin pulvinar purus, 85 eget luctus lorem tincidunt vitae. Vestibulum eu mattis neque. Nulla 86 in tortor id urna dapibus gravida a vel leo. 79 The Pan-STARRS Image Processing Pipeline performs the processing 80 needed to downloaded, archive, and process all images obtained by the 81 Pan-STARRS telescopes. This article describes the overall data 82 analysis system. 87 83 88 84 \end{abstract} … … 91 87 \keywords{Surveys:\PSONE } 92 88 93 % \section{INTRODUCTION}\label{sec:intro} 89 \section{INTRODUCTION}\label{sec:intro} 90 91 This is the second in a series of seven papers describing the 92 Pan-STARRS1 Surveys, the data reduction techiques and the resulting 93 data products. This paper (Paper II) describes how the various data 94 processing stages are organised and implemented in the Imaging 95 Processing Pipeline (IPP), including details of the the processing 96 database which is a critical element in the IPP infrastructure . 97 98 %Chambers et al. 2017 (Paper I) 99 %The Pan-STARRS\,1 Surveys 100 \citet[][Paper I]{chambers2017} 101 provides an overview of the Pan-STARRS System, the design and 102 execution of the Surveys, the resulting image and catalog data 103 products, a discussion of the overall data quality and basic 104 characteristics, and a brief summary of important results. 105 106 %Magnier et al. 2017 (Paper II) 107 %Pan-STARRS Data Processing Stages 108 %\citet[][Paper II]{magnier2017c} 109 %describes how the various data processing stages are organised and implemented 110 %in the Imaging Processing Pipeline (IPP), including details of the 111 %the processing database which is a critical element in the IPP infrastructure . 112 113 %Waters et al. 2017 (Paper III) 114 %Pan-STARRS Pixel Processing : Detrending, Warping, Stacking 115 \citet[][Paper III]{waters2017} 116 describes the details of the pixel processing algorithms, including detrending, warping, and adding (to create stacked images) and subtracting (to create difference images) and resulting image products and their properties. 117 118 119 %Magnier et al. 2017 (Paper IV) 120 %Pan-STARRS Pixel Analysis : Source Detection 121 \citet[][Paper IV]{magnier2017a} 122 describes the details of the source detection and photometry, including point-spread-function and extended source fitting models, and the techniques for ``forced" photometry measurements. 123 124 %Magnier et al. 2017 (Paper V) 125 %Pan-STARRS Photometric and Astrometric Calibration 126 \citet[][Paper V]{magnier2017b} 127 describes the final calibration process, and the resulting photometric and astrometric quality. 128 129 130 %Flewelling et al. 2017 (Paper VI) 131 %Pan-STARRS 1 Database and Data Products 132 \citet[][Paper VI]{flewelling2017} 133 describes the details of the resulting catalog data and its organization in the Pan-STARRS database. 134 % 135 % 136 \citet[][Paper VII]{huber2017} 137 %Huber et al. 2017 (Paper VII) 138 describes the Medium Deep Survey in detail, including the unique issues and data products specific to that survey. The Medium Deep Survey is not part of Data Release 1. (DR1) 139 140 % 141 The Pan-STARRS1 filters and photometric system have already been 142 described in detail in \cite{2012ApJ...750...99T}. 143 144 {\color{red} {\em Note: These papers are being placed on arXiv.org to 145 provide crucial support information at the time of the public 146 release of Data Release 1 (DR1). We expect the arXiv versions to 147 be updated prior to submission to the Astrophysical Journal in 148 January 2017. Feedback and suggestions for additional information 149 from early users of the data products are welcome during the 150 submission and refereeing process.}} 94 151 95 152 \section{IPP Software Subsystems} … … 852 909 \subsection{UH Cray Cluster} 853 910 911 \acknowledgments 912 913 The Pan-STARRS1 Surveys (PS1) have been made possible through 914 contributions of the Institute for Astronomy, the University of 915 Hawaii, the Pan-STARRS Project Office, the Max-Planck Society and its 916 participating institutes, the Max Planck Institute for Astronomy, 917 Heidelberg and the Max Planck Institute for Extraterrestrial Physics, 918 Garching, The Johns Hopkins University, Durham University, the 919 University of Edinburgh, Queen's University Belfast, the 920 Harvard-Smithsonian Center for Astrophysics, the Las Cumbres 921 Observatory Global Telescope Network Incorporated, the National 922 Central University of Taiwan, the Space Telescope Science Institute, 923 the National Aeronautics and Space Administration under Grant 924 No. NNX08AR22G issued through the Planetary Science Division of the 925 NASA Science Mission Directorate, the National Science Foundation 926 under Grant No. AST-1238877, the University of Maryland, and Eotvos 927 Lorand University (ELTE) and the Los Alamos National Laboratory. 928 929 \bibliographystyle{apj} 930 % \bibliography{lib}{} 931 \input{datasystem.bbl} 932 854 933 \end{document} 855 934
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