# The VST Photometric H$\alpha$ Survey of the Southern Galactic Plane and
Bulge (VPHAS+)

J. E. Drew1, E. Gonzalez-Solares2, R. Greimel3, M. J. Irwin2, A. Kupcu
Yoldas2, J. Lewis2, G. Barentsen1, J. Eislöffel4, H. J. Farnhill1, W. E.
Martin1, J. R. Walsh5, N. A. Walton2, M. Mohr-Smith1, R. Raddi6, S. E. Sale7,
N. J. Wright1, P. Groot8, M. J. Barlow9, R. L. M. Corradi10, J. J. Drake11, J.
Fabregat12, D. J. Frew13, B. T. Gänsicke6, C. Knigge14, A. Mampaso10, R. A. H.
Morris15, T. Naylor16, Q. A. Parker13, S. Phillipps14, C. Ruhland1, D.
Steeghs6, Y.C. Unruh17, J. S. Vink18, R. Wesson19, A. A. Zijlstra20
1School of Physics, Astronomy & Mathematics, University of Hertfordshire,
College Lane, Hatfield, Hertfordshire, AL10 9AB, U.K.
2Institute of Astronomy, Cambridge University, Madingley Road, Cambridge, CB3
OHA, U.K.
3IGAM, Institute of Physics, University of Graz, Universitätsplatz 5, Graz,
Austria
4Thüringer Landessternwarte, Sternwarte 5, 07778, Tautenburg, Germany
5ESO Headquarters, Karl-Schwarzschild-Strasse 2, 85748 Garching, Germany
6Department of Physics, University of Warwick, Gibbet Hill Road, Coventry, CV4
7AL, U.K¿
7Rudolf Peierls Centre for Theoretical Physics, Keble Road, Oxford, OX1 3NP
8Afdeling Sterrenkunde, Radboud Universiteit Nijmegen, Faculteit NWI, Postbus
9010, 6500 GL Nijmegen, The Netherlands
9University College London, Department of Physics & Astronomy, Gower Street,
London WC1E 6BT, U.K.
10 Instituto de Astrofisica de Canarias, 38200 La Laguna, Tenerife, Spain
11Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA
02138, U.S.A.
12Observatorio Astrónomico, Universidad de Valencia, Catedrático José Beltrán
2, 46980 Paterna, Spain
13Department of Physics & Astronomy, Macquarie University, NSW 2109, Australia
14School of Physics & Astronomy, University of Southampton, Southampton, SO17
1BJ, U.K.
15School of Physics, Bristol University, Tyndall Avenue, Bristol, BS8 1TL,
U.K.
16School of Physics, University of Exeter, Stocker Road, Exeter, EX4 4QL, U.K.
17Department of Physics, Blackett Laboratory, Imperial College London, Prince
Consort Road, London, SW7 2AZ, U.K.
18Armagh Observatory, College Hill, Armagh, Northern Ireland, BT61 9DG, U.K.
19European Southern Observatory, Alonso de Córdova 3107, Casilla 19001,
Santiago, Chile
20Jodrell Bank Centre for Astrophysics, School of Physics & Astronomy,
University of Manchester, Oxford Road, Manchester M13 9PL, U.K.

###### Abstract

The VST Photometric H$\alpha$ Survey of the Southern Galactic Plane and Bulge
(VPHAS$+$) is surveying the southern Milky Way in $u,g,r,i$ and H$\alpha$ at
$\sim$1 arcsec angular resolution. Its footprint spans the Galactic latitude
range $-5^{\rm o}<b<+5^{\rm o}$ at all longitudes south of the celestial
equator. Extensions around the Galactic Centre to Galactic latitudes $\pm
10^{\circ}$ bring in much of the Galactic Bulge. This ESO public survey, begun
on 28th December 2011, reaches down to $\sim$20th magnitude (10$\sigma$) and
will provide single-epoch digital optical photometry for $\sim$300 million
stars. The observing strategy and data pipelining is described, and an
appraisal of the segmented narrowband $H\alpha$ filter in use is presented.
Using model atmospheres and library spectra, we compute main-sequence $(u-g)$,
$(g-r)$, $(r-i)$ and $(r-H\alpha)$ stellar colours in the Vega system. We
report on a preliminary validation of the photometry using test data obtained
from two pointings overlapping the Sloan Digital Sky Survey. An example of the
$(u-g,g-r)$ and $(r-H\alpha,r-i)$ diagrams for a full VPHAS+ survey field is
given. Attention is drawn to the opportunities for studies of compact nebulae
and nebular morphologies that arise from the image quality being achieved. The
value of the $u$ band as the means to identify planetary-nebula central stars
is demonstrated by the discovery of the central star of NGC 2899 in survey
data. Thanks to its excellent imaging performance, the VST/OmegaCam
combination used by this survey is a perfect vehicle for automated searches
for reddened early-type stars, and will allow the discovery and analysis of
compact binaries, white dwarfs and transient sources.

###### keywords:

surveys – stars: emission line – Galaxy: stellar content

##  1 Introduction

The $H\alpha$ emission line is well-known as a tracer of diffuse ionized
nebulae and as a marker of pre- or post-main sequence status among spatially-
unresolved stellar sources. Since these objects – both nebulae and stars –
represent relatively short-lived phases of evolution, they amount to a
minority population in a mature galaxy like our own. Their relative scarcity
has in the past stood in the way of developing and testing models for these
crucial evolutionary stages.

In the southern hemisphere, the search for planetary nebulae (PNe) has been
served well by H$\alpha$ imaging surveys carried out by the UK Schmidt
Telescope (Parker et al 2005, 2006 and other more recent works). Nevertheless,
VPHAS+ will have a decisive impact on studies of complex or smaller nebulae of
all types, ranging from optically-detectable ultra-compact and compact HII
regions, to nebulae around YSOs (including associated jets and HH objects),
through PNe, to extended emission from D-type symbiotic stars and supernova
remnants. The superb spatial resolution, dynamic range, and likely photometric
accuracy of the VPHAS+ images warrant a step forward in our knowledge of the
population and detailed characteristics of these object classes.

For southern point sources with emission the situation is very different:
there has been little updating of the available catalogues since the work of
Stephenson & Sanduleak (1971) that was limited to a depth of 12th magnitude.
The major groups of emission line stars that remain as challenges to our
understanding include all types of massive star (O stars, supergiants,
luminous blue variables, Wolf-Rayet stars, various types of Be star), post-AGB
stars, pre-main-sequence stars at all masses, active stars and compact
interacting binaries. Within the disc of the Milky Way, the available samples
of these objects are typically modest and heterogeneous. Fixing this deficit
via a uniform search of the Galactic Plane for these rare object classes
motivated the photometric H$\alpha$ survey of the southern Galactic Plane,
first proposed for the VLT Survey Telescope (VST) in 2004. This paper
describes the realisation of this ESO public survey, now known as the VST
Photometric H$\alpha$ Survey of the Southern Galactic Plane and Bulge
(VPHAS+).

When first proposed, VPHAS (without the plus sign) was envisaged as the
counterpart to the INT/WFC Photometric H$\alpha$ Survey of the Northern
Galactic Plane (IPHAS, Drew et al 2005), that had begun in August 2003. IPHAS
is a digital imaging survey made up of Sloan $r$, $i$ and narrowband $H\alpha$
exposures, reaching to $\sim$20th magnitude, that takes in all Galactic
longitudes north of the celestial equator in the latitude range
$-5^{\circ}<b<+5^{\circ}$. This is all but complete with the new release of a
catalogue of $\sim$200 million unique objects drawn from 93 percent of the
survey’s footprint (Barentsen et al, 2014). During the initial VST public
survey review process, it was agreed that VPHAS should broaden in scope to
also incorporate the Sloan $u$ and $g$ bands (proposed for a separate survey),
that are particularly useful in picking out OB stars, white dwarfs and other
blue-excess objects. With this upgrade to 5 bands, the renamed VPHAS$+$ became
an all-purpose digital optical survey of the southern Galactic Plane, capable
of delivering data at a spatial resolution of $\sim$1 arcsec or better. As
well as fulfilling the role of southern counterpart of IPHAS, VPHAS$+$ is also
the counterpart to UVEX, the UV-Excess Survey of the Northern Galactic Plane
(Groot et al, 2009) that, at the time of writing, continues on the Isaac
Newton Telescope in La Palma.

The final augmentation of the VPHAS$+$ survey footprint came in 2010 on
expanding its footprint to match that of the similarly high spatial-resolution
near-infrared survey, VISTA Variables in the Via Lactea (VVV, Minniti et al
2011). The survey footprint now includes the Galactic Bulge to a latitude of
$|b|<10^{\circ}$, across the longitude range $-10^{\circ}<\ell<+10^{\circ}$.
The VVV $z$, $Y$, $J$, $H$ and $K_{s}$ survey of much of the Bulge and inner
Galactic disc is already complete.

VPHAS$+$ is poised to become the homogeneous digital optical imaging survey of
the Galactic Plane and Bulge, at $\sim$1 arcsec angular resolution, that will
provide a uniform database of stellar spectral energy distributions, from
which a range of colour-magnitude and colour-colour diagrams of well-
established utility can be derived. Once calibrated to the expected precision
of 2 to 3 percent, like its northern counterparts, VPHAS$+$ will be
quantitatively far superior to the photographic surveys of the last century,
and will offer significant added value in the form of calibrated narrowband
H$\alpha$ data. Some of the science enabled is illustrated by the studies that
the northern surveys, IPHAS and UVEX, have already stimulated. These have
included a number of works reporting the discovery of emission line stars,
ranging from young objects (e.g. Valdivieso et al, 2009, Vink et al 2008,
Barentsen et al 2011, Raddi et al 2013) to evolved object classes such as
symbiotic and cataclysmic binaries and compact planetary nebulae (e.g. Corradi
et al 2010, Witham et al 2006, Viironen et al 2009 and Wesson et al 2008). The
diagnostic power to be expected from the blue $u$ and $g$ bands has been
appraised in a series of papers presenting UVEX and early follow-up
spectroscopy (Groot et al 2009, Verbeek et al 2012a and 2012b).

Figure 1: The VPHAS$+$ survey footprint plotted in Galactic coordinates. All
2269 fields are shown in outline. Different colours, as specified in the key,
are used to identify the observations obtained for each field by 1st January
2014 – essentially 2 years after the start of data-taking.

Any comprehensive survey of the Galactic Plane clearly targets the main mass
component of our own Galaxy, made up of stars, gas and dust. When a spatial
resolution of 1 arcsec is combined with wide area coverage spanning 100s of
square degrees, as in IPHAS, UVEX and VPHAS$+$, it becomes possible to exploit
the stellar photometry adaptively to solve for the distribution of both the
stars and dust making up the optically-accessible Galactic disk - by means
that are similar to those already attempted, based on the near-infrared 2MASS
survey (Drimmel & Spergel 2001, Marshall et al 2006). Methods to achieve
3-dimensional mapping of this kind, now incorporating the direct sensitivity
of the H$\alpha$ narrowband to stellar intrinsic colour, are starting to take
shape (Sale et al 2009; Sale 2012). This development coincides with the
approach of the operations phase of Europe’s next major astrometric mission,
Gaia. Indeed, the rich dynamical picture that Gaia will build of the Milky Way
over the next decade will be very effectively complemented by stellar energy
distributions measured for millions of stars from the current generation of
ground-based optical, and near-infrared, wide-field surveys. VPHAS$+$, with
its haul of photometry in 5 optical bands on $\sim$300 million objects, is set
to take its place as one of them.

This paper presents the main features of VPHAS$+$, including a description of
its execution, the data processing and the nature of the photometric colour
information it provides. We begin in the next section with a presentation of
the observing strategy and the data reduction techniques in use. We then turn
to a description and evaluation of the narrowband H$\alpha$ filter procured
for this survey, in section 3. Following this, in section 4, we present
tailored synthetic photometry of main sequence and giant stars that provides
insights into the photometric diagrams that may be generated from survey data.
An exercise in photometric validation is described in Section 5 in which Sloan
Digital Sky Survey (SDSS) data are compared with VST observations. The scene
is then set for an example of VPHAS+ photometry extracted across the entire
square-degree footprint of a single survey field (Section 6). In Sections 7
and 8, we outline the applications of VPHAS$+$ to spatially-resolved nebular
astrophysics and in the time domain. The paper ends in Section 9 with a
summary, examples of early survey exploitation, and a forward look to the
first major data release.

##  2 Survey observations and data processing

###  2.1 VPHAS$+$ specification

The footprint of the survey is shown in figure 1. The OmegaCAM imager (Kuijken
2011) on the VST provides a field size of a full square degree, captured on a
4 x 8 CCD mosaic. After allowing for some modest overlap between adjacent
fields, we arrived at a set of 2269 field centres that will cover the desired
Galactic latitude band $-5^{\circ}<b<+5^{\circ}$ at all southern-hemisphere
Galactic longitudes, as well as incorporate the Galactic Bulge extensions to
$-10^{\circ}<b<+10^{\circ}$ near the Galactic Centre. The survey footprint
extends across the celestial equator by a degree or two to achieve an overlap
with the northern hemisphere surveys IPHAS and UVEX of $\sim$100 sq.deg.
altogether. This is to create the opportunity for some direct photometric
cross-calibration.

The target depth of the survey is to reach to at least $\sim$20th magnitude,
at 10$\sigma$, in each of the Sloan $u$, $g$, $r$ and $i$ broadband filters
and narrowband H$\alpha$. The bright limit consistent with this goal is
typically 12–13th magnitude. Presently, all VPHAS+ photometric magnitudes are
expressed in the Vega system. The original concept was to collect the data in
all 5 bands contemporaneously, in order to build a uniform library of snapshot
photometric spectral energy distributions for 200 million or more stars.
Practical constraints have modified this to the extent that the blue filters
($u$, $g$) are observed separately from the reddest ($i$ and $H\alpha$), with
the $r$ band serving as a linking reference that is observed both with $u$,$g$
and $i$,$H\alpha$. The aim is also to keep the spatial resolution close to 1
arcsecond. OmegaCAM and the Paranal site are well suited to this in that the
camera pixel size is 0.21 arcsec, projected on sky, and the median seeing
achieved is better than 1 arcsec (on occasion falling to as little as 0.6
arcsec).

As a means to obtaining better quality control, and to ensure that only a
minimal fraction of the survey footprint is missed due to the pattern of gaps
between the CCDs in the camera mosaic, every field is imaged at 2 or 3 offset
pointings. This strategy has been carried over from IPHAS and UVEX in the
northern hemisphere, and has the consequence that the majority of imaged
objects will be detected twice some minutes apart. In the $r$ band, there will
be two arbitrarily-separated epochs of data, with typically two detections at
each of them (i.e. 4 altogether).

###  2.2 VPHAS$+$ observations

The VST is a service-observing facility, with all programmes queued for
execution as and when the ambient conditions meet programme requirements.
VPHAS+ survey field acquisition began on 28th December 2011. Normally the
constraint set includes a seeing upper bound of 1.2 arcsec: this is only set
at a lower, more stringent value for fields expected to present a particularly
high density of sources (e.g. in the southern Bulge). In order that the seeing
achieved in the $u$ band is not greatly different from that in $i$ at the
opposite end of the optical range, it is advantageous to separate acquisition
of blue data from red – hence a split between ’blue’ ($u$,$g$,$r$) and ’red’
($H\alpha$,$r$,$i$) observing blocks has been implemented. This split also
permits the use of different moon distance and phase constraints, such that
blue data are obtained when the moon is less than half full at an angular
separation of not less than 60 degrees, while the limits for red data are set
at 0.7 moon illumination and a minimum angle of 50 degrees. Avoiding bright-
moon conditions is important in order to limit the amount of moonlight mixed
in with diffuse H$\alpha$ emission in the reduced images. No requirement has
been placed on the time elapsing between acquisition of blue and red data.
However, the more forgiving constraints on the acquisition of the latter has
meant that these are typically executed sooner than the former, with the
result that many more fields have red data already than have blue (see fig 1).
In all cases, the final constraint is that the sky is required to be clear, if
not necessarily fully photometric.

An impressive feature of the camera, OmegaCAM, is its potential to deliver
remarkably undistorted point-source images all the way across the 1-degree
field of view. To realise this, it is critical that the VST has an actively-
controlled primary. The operational price for this, at the present time, is
that image analysis and correction has to be carried out at every filter
change or after longer slews. The overhead added by this is about 3 minutes.
To reduce the impact of this, observations of sets of 3 neighbouring fields
are scheduled together, so that image analysis need only take place every
15-30 minutes – not much more often than would be essential, in any case, to
compensate for the telescope’s tracking movement. As a result,
’contemporaneous’ in the context of VPHAS+ data-taking means that all 3 blue,
or red, filters are typically exposed within 40-50 minutes of each other (cf.
IPHAS, where the more compact camera allows much faster operation, bringing
this elapsed time down to under 10 minutes). However, the time difference
between the blue and red observing blocks for a given field, i.e. the
$u$/$g$/$r$ data collection and $H\alpha$/$r$/$i$ data collection, can be
anything from a few hours to more than a year.

Figure 2: The transmission profiles of the Sloan $u$, $g$, $r$ and $i$, and
narrowband $H\alpha$ filters used in all VPHAS$+$ observations. The $r$ and
$H\alpha$ profiles are shown as a dashed line and in red respectively just to
clearly distinguish them from each other and the $i$ band. Each profile has
been multiplied by the CCD response function and a model of atmospheric
throughput (Patat et al 2011). The very-nearly grey losses due to the
telescope optics (a further scaling of approximately 0.6) have not been folded
in.  Figure 3: An illustration of the VPHAS$+$ offset pattern as it applies to
the segmented H$\alpha$ filter with extra vignetting due to the T bars
separating the 4 segments. The first pointing is to lower right, as drawn – a
conservative estimate of the exposed unvignetted area is shown in black. The
exposed/unvignetted areas of the second and third pointings are shown in green
and blue respectively. The vertical and horzontal scales are numbered in
pixels (RA increasing to the right, declination increasing upwards).

We provide a reminder of the passbands of the Sloan broadband filters in use,
along with that of the narrowband H$\alpha$, in fig 2. They are shown scaled
by a typical CCD response function and a model of the atmospheric transmission
(Patat et al 2011). The exposure times used for the different filters, the
number of exposures in each and the median seeing achieved, up to December
2013, are set out in Table 1. From early on, during the commissioning phase of
the telescope, it became clear that tracking is usually good enough that even
the 150 sec $u$ exposures, our longest, do not have to be guided in most
circumstances. Indeed, experience is showing that it is safer to rely on the
tracking, rather than the autoguider, to maintain good image quality in the
most dense star fields.

Table 1: Observations obtained per survey field. The median seeing quoted is derived from data in hand by December 2013. Filter | Exposure | No. of | median seeing
---|---|---|---
| time (secs) | offsets | (arcsec)
Blue observation blocks
$u$ | 150 | 2 | 1.01
$g$ | 40111Up to 19th February 2013, $g$ exposure times were 30 sec. | 3 | 0.88
$r$ | 25 | 2 | 0.80
Red observation blocks
$H\alpha$ | 120 | 3 | 0.84
$r$ | 25 | 2 | 0.82
$i$ | 25 | 2 | 0.77

The pattern of offsets used for each field is illustrated in fig 3. The shifts
are relatively large, with the outer pointings differing by $-$588 arcsec in
the RA direction and $+$660 arcsec in declination. The choices made have
largely been driven by the characteristics of the narrowband H$\alpha$ filter
(discussed below in section 3), but they also convey the advantage of greatly
increasing the overlaps between neighbouring fields. Just the two outermost
pointings are used when exposing the $u$, $r$ and $i$ filters. This leaves
0.4% of the survey footprint unexposed. This changes to complete coverage on
including the third intermediate offset, as is the policy for the $H\alpha$
and $g$ filters.

In accordance with ESO’s standard procedures, data are evaluated soon after
collection by Paranal staff and graded before transfer to the archive in
Garching and to the Cambridge Astronomy Survey Unit (CASU) in Cambridge. If
the applied constraints are significantly violated, the observation block is
returned to the queue.

###  2.3 Data pipeline

####  2.3.1 Initial Processing

Figure 4: Flowchart identifying the main VST data processing steps. At the
present time all steps up to the reduced single-band catalogue are undertaken
at CASU. Band merging is performed at the University of Hertfordshire.

From February 29th 2012 raw VST data have been routinely transferred from
Paranal to Garching over the Internet. For each observation the imaging data
are stored in a Multi-Extension FITS file (MEF) with a primary header
describing the overall characteristics of the observation (pointing, filter,
exposure time, etc.) and thirty two image extensions, corresponding to each of
the CCD detectors, with further detector-level information in the secondary
headers. The 32-bit integer raw data files are Rice-compressed at source using
lossless compression (e.g. Sabbey 1998). The files are then checked and
ingested into the ESO raw data archive in Garching. As soon as the data for
any given night become available they are automatically transferred to the
Cambridge Astronomy Survey Unit (CASU) for further checks and subsequent
processing. The VST web pages at CASU provide an external interface for both
monitoring processing status (http://casu.ast.cam.ac.uk/vst/data-processing/)
and overall survey progress and access (http://casu.ast.cam.ac.uk/vstsp/).

The processing sequence is similar to that used for the IPHAS survey of the
northern Galactic Plane (e.g. Gonzalez-Solares et al 2008), while the higher
level control software is based on that developed for the VISTA Data Flow
System (VDFS, Irwin et al 2004). Here we briefly outline the processing steps
illustrated in figure 4, emphasising the main differences relative to the
current VDFS standard. A more detailed description of the VST processing
pipeline is currently in preparation (Yoldas et al 2014).

Science images are first debiassed. Full two-dimensional bias removal is
necessary due to amplifier glow during readout being present in some
detectors. The master bias frames are updated daily from calibration files
taken as part of the operational cycle. The OmegaCAM detectors are linear to
better than 1% over their usable dynamic range removing the need for a
linearity correction. Hence this stage in the pipeline processing (figure 4),
although part of the pipeline architecture, is currently bypassed.

Flatfield images in each band are constructed by combining a series of
twilight sky flats obtained in bright sky conditions. The timescale to obtain
sequences of these for all deployed filters is typically one to two weeks. So
as to adequately trace the variations in the pattern and level of scattered
light in these flats, the master flats derived from them are updated on a
monthly cycle (how these are corrected for scattered light is described in
Section 2.3.3). Four of the detectors, in extensions 29–32, suffer from inter-
detector cross-talk, whereby saturated bright stars in one detector can cause
noticeable positive or negative low level ($\approx$0.1%) ghost objects in
adjacent detectors. The impact of these is minimised in the pipeline by
applying a pre-tabulated cross-talk correction matrix to each of the affected
images.

The flatfield sequences plus bad pixel masks are used to generate the
confidence (weight) maps (e.g. Irwin et al 2004) used later during catalogue
generation and any subsequent image stacking or large area mosaicing. After
flatfielding, science images generally have well behaved sky backgrounds which
makes subsequent image processing straightforward. Where direct scattered
light is present in them, it is an additive phenomenon that is dealt with
automatically during object catalogue generation.

Figure 5: The left-hand panel shows an example of the deduced scattered light
component present in June 2013 $r$-band data. The right-hand panel shows the
corrected outcome. These maps were constructed using 286318 APASS object
matches over the square-degree field. Before correction scattered light
gradients amounting to a 20 to 25 percent variation from corners to centre are
present. This flattens to around $\pm$2 percent.

The redder passbands used in VST observations, in this case the $i$-band, show
fringing patterns at $\approx$2% of the sky background level. Defringing is
done using a standard CASU procedure (Irwin & Lewis 2001). The fringe frames
used are derived from other VST public survey data taken as close as possible
in time at higher Galactic latitude. This approach works because the fringe
pattern induced by sky emission lines at Paranal is quite stable over long
periods. The fringe frames are automatically scaled and subtracted from each
science image reducing the residual fringing level to well below the sky
noise.

Catalogue generation is based on IMCORE222Software publicly available from
http://casu.ast.cam.ac.uk (Irwin, 1985) and makes direct use of the confidence
maps, derived from the flat fields, to suitably weight down unreliable parts
of the images. This step includes object detection, parameterization and
morphological classification, together with generation of a range of quality
control information. Because of the extensive presence of diffuse emission
throughout the southern Galactic Plane, particularly in H$\alpha$, a version
of each affected image is cleaned of nebulosity using the NEBULISER333Software
publicly available from http://casu.ast.cam.ac.uk, see also Irwin (2010) for
the purpose of catalogue generation only. This achieves a more careful removal
of background and ultimately leads to more complete and, on average, more
faithful object detection than in the absence of this step.

With object catalogues available for every VPHAS+ survey image, it is then
possible to improve the rough World Coordinate System (WCS) based on the
telescope pointing and general system characteristics. The WCS is
progressively refined using matches between detected objects and the 2MASS
catalogue (Skrutskie et al 2006). Despite the large field of view, the VST
focal plane is almost free of distortion, and a standard tangent plane
projection yields residual systematics of $\sim$25mas over the entire field.

####  2.3.2 Photometric Calibration

Provisional photometric calibration is based on a series of standard star
fields observed each night (e.g. Landolt 1992). For each night a zeropoint and
error estimate using the observations of all the standard fields in each
filter is derived. The flatfielding stage nominally places all detectors on a
common internal gain system implying, in principle, that a single zeropoint
suffices to characterise the whole focal surface. Colour equations are used to
transform between the passbands in use on the VST and the Johnson-Cousins
system of the published standard-star photometry. The calibration is currently
in a VST system that uses the SED of Vega as the zero-colour, almost zero-
magnitude, reference object.

The $u$ band data are the most challenging to calibrate. As this part of
Vega’s spectrum, and also the average standard star plus the detector reponse,
are falling rapidly it would be surprising if there were no offsets in $u$ due
to nonlinearities in the required colour transforms and, perhaps, to
degenerate colour transforms for hotter stars. Early experience of working
with $u$ data do indeed suggest that offsets of up to a few tenths of a
magnitude are sometimes present (see Sections 5 and 6).

The colour transforms currently in use to define the VPHAS internal system are
given below.

$\displaystyle u_{VST}$ $\displaystyle=$ $\displaystyle U+0.035\,(U-B)$
$\displaystyle g_{VST}$ $\displaystyle=$ $\displaystyle B-0.405\,(B-V)$
$\displaystyle r_{VST}$ $\displaystyle=$ $\displaystyle R+0.255\,(V-R)$
$\displaystyle i_{VST}$ $\displaystyle=$ $\displaystyle I+0.215\,(R-I)$
$\displaystyle H\alpha_{\,VST}$ $\displaystyle=$ $\displaystyle
R+0.025\,(V-R)$

The transform for the narrowband H$\alpha$ is an approximate initial solution
needed for the subsequent illumination correction stage. At catalogue
bandmerging this is superceded (see Section 4).

Figure 6: Photometric errors in VPHAS$+$ data as a function of magnitude. All
data are drawn from an area of $\sim$0.2 sq.deg in field 1679, positioned
$\sim$20 arcmin E of Westerlund 2. The left hand panel refers to blue
observing block (OB) data, while the right refers to red OB data. The coloured
histogram in each component plot shows the mean absolute deviation of the
magnitude difference, $(m_{1}-m_{2})$, per 0.25-magnitude bin, while the bin
means of the suitably-corrected pipeline estimate of the random error on this
difference are in black. The 10- and 5-$\sigma$ magnitude limits are specified
in the brackets next to each filter name. The rising observed mean deviation
seen in $g$ at the bright end is due to the onset of saturation.

####  2.3.3 Illumination Correction

The main difficulty in deriving an accurate photometric calibration over the
one degree field arises from the multiplicative systematics caused by
scattered light in the flatfields. The VST (at least up to the introduction of
baffles early in 2014) has proved particularly susceptible to variable
scattered light. Its impact has varied from month to month depending on
conditions prevalent at the time the flat-field sequences were taken. An
illustration of the amount and character of the master-flat correction
required is provided in figure 5.

The scattered light is made up of multiple components with different
symmetries and scales. These range from $\approx$10 arcsec with x-y
rectangular symmetry, e.g. due to scattering off masking strips above the CCD
readout edges, to large fractions of the field due to radial concentration of
light in the optics and to non-astronomical scattered light entering obliquely
in flatfield frames. After some experimentation, and external verification, we
found that the APASS all-sky photometric g,r,i catalogues
(http://www.aavso.org/apass) provide a reliable working solution to the
illumination correction problem inherent in VST data (see fig 5). These
catalogues also provide an independent overall photometric calibration tied to
the SDSS AB magnitude system and will be used in future updates to define an
alternative finer-grained temporal AB magnitude zeropoint.

All filters used are treated in the same manner with colour equations set up
to define transformations between the APASS g,r,i SDSS-like calibration and
the VPHAS+ u,g,r,i,Hα internal system.

Illumination corrections are re-derived for each filter once a month.
Application of these corrections via the master flats reduces the residual
systematics across the entire field to below the 1% level for the broadband
filters and to within 2% for the segmented narrowband H$\alpha$, except in
vignetted regions (see Section 3).

####  2.3.4 Quality Control

In addition to the usual VDFS quality-control monitoring of average stellar
seeing and ellipticity, sky surface brightness and noise properties, we have
also initiated a more detailed analysis of the image properties based on
inter-detector comparisons.

The well-aligned coplanar detector array coupled with the curved focal surface
is extremely sensitive to imperfections in focus which are relatively easy to
detect using the detector-level average seeing measurement variation available
for each of the 32 detectors. Likewise the variation in average stellar
ellipticity from each detector over the field is used to monitor rotator angle
tracking problems.

All of this information can be used in addition to the observation block (OB)
grades provided by ESO and is incorporated within all data product files and
also the progress database.

###  2.4 Limiting magnitudes and errors

The present convention for VPHAS+ and this paper is that all magnitudes are
expressed in the Vega system, which imposes zero intrinisc colour for A0
stars. The 5-sigma limiting magnitudes commonly achieved per exposure range
from 20.5–21.0 for H$\alpha$ up to 22.2–22.7 for Sloan $g$. The 10$\sigma$
limits are about 1 magnitude brighter.

Every source flux or magnitude determined via the pipeline has a formal error
associated with it. We provide an example of how these compare with empirical
magnitude differences, by extracting a sample of stars from a 0.25 sq.deg
catalogue, cut out from the survey field including Westerlund 2 (field 1679,
see also sec 6) in order to examine the pattern of errors (fig 6). The sky
area chosen is offset from the cluster to the east by $\sim$20 arcmin and
exhibits moderate diffuse ionised nebulosity. In the southern Plane, the
presence of some nebulosity, particularly affecting $r$ and $H\alpha$
exposures, is more the rule than the exception. Sources classified as probable
stars in both $g$ and $r$ (blue filter set, left panel) and $r$ and $i$ (red
filter set, right hand panel) in two consecutive offset exposures have been
selected. The selection also required that each extracted magnitude was
unaffected by vignetting and bad pixels (confidence level $>95$). This step is
particularly important for H$\alpha$ given the extra vignetting introduced by
the cross bars of the segmented filter (see Section 3).

The faintest stars that might have been included in the plots for $i$ and blue
$r$ (top row in fig 6) are absent because of a requirement that every included
source should also be picked up in, respectively, red $r$ and $g$. Fainter
objects than the apparent limits certainly exist in these bands. This feature
follows directly from the typically red colours of Galactic Plane stars at
magnitudes fainter than $\sim$13 that are the target of this survey. For the
same reason, it is not uncommon for the $u$ band source counts to be one or
more orders of magnitude lower than those of the $i$ band. The role of the $u$
band is to pick out the unusual rather than to characterise the routine.

To bring out the systematic effects present, the specific comparison made in
fig 6 is between the bin means of the absolute magnitude differences,
$|m_{1}-m_{2}-\delta|$, between the two exposures, and the expected random
error on the difference derived from the pipeline rms errors on the individual
magnitude measurements. The quantity, $\delta$, is the median magnitude
difference computed from all bright stars down to 18th magnitude ($r$, $i$ and
H$\alpha$), or 19th magnitude ($u$ and $g$). This was small in all cases – the
largest value being 0.011 for $u$. On the other hand, the correction applied
to the pipeline errors was, first, to multiply the single-measurement
magnitude error by $\sqrt{2}$ to give the rms error on the $(m_{1}-m_{2})$
difference, and then to multiply by $\sqrt{2/\pi}$ in order to convert the
measure of dispersion from rms to a mean deviation.

At magnitudes brighter than 18–19 in fig 6, the scatter in the empirical
results can be seen to be appreciably greater than that ’predicted’ for the
random component by the pipeline. The scale of the difference indicates that a
further error component of $\sim$0.01–0.02 magnitudes is present. The amount
and filter-dependence of these levels of error are entirely consistent with
the uncertainties estimated above for the flatfield and illumination
corrections: as noted in section 2.3.3, the VST is presently prone to quite
high and variable levels of scattered light. Practical remedies for this are
under consideration by ESO – when implemented these should tighten up the
error budget.

The enhanced mean magnitude difference seen for $g<13$ in fig 6 is typical of
what is seen as saturation effects begin to set in. For all the other bands,
in this example, saturation sets in at magnitudes a little brighter than 12th.
A safe working assumption across VPHAS+ would be that saturation is never
troublesome at magnitudes fainter than 13, but always an issue for magnitudes
brighter than 12.

Figure 6 identifies the 5$\sigma$ and 10$\sigma$ magnitude limits for each
filter achieved in this representative example. The seeing at the time of
these observations, as measured in the pipeline, ranged from 0.8 to 1 arcsec
(cf Table 1).

##  3 The narrowband $H\alpha$ filter

###  3.1 Overview

Table 2: Summary of filter segment properties Segment | sky quadrant | CCDs covered | centre, mean, corner CWL | mean integrated
---|---|---|---|---
|  |  | (Å) | throughput (Å)
A | SW | 1 – 8 | 6580.2 – 6585.4 – 6595.3 | 98.64
B | SE | 17 – 24 | 6596.1 – 6585.4 – 6578.9 | 103.38
C | NE | 25 – 32 | 6582.8 – 6591.9 – 6599.8 | 99.74
D | NW | 9 – 16 | 6581.7 – 6594.3 – 6603.8 | 99.49
Figure 7: The segmented H$\alpha$ filter, photographed in the lab soon after
receipt and just prior to measuring its transmission. The filters to either
side of the 2$\times$2 array of H$\alpha$ segments, transmit as $r$-band and
cover the guide CCDs.

A filter required to select a narrow band across a large 27$\times$27 cm2
image plane is a challenging fabrication problem. At the telescope, the filter
in use for VPHAS$+$ is known as NB-659. At the time it was commissioned in
2006, the purchase of a single-piece narrowband filter was offered only by one
supplier and was well beyond budget. This left the 4-segment option as the
achievable alternative.

The H$\alpha$ filter was constructed based on a specification supplied by the
OmegaCAM consortium, setting as goal a central wavelength of 6588 Å, and a
bandpass of 107 Å. It was delivered in the summer of 2009, and was shortly
thereafter tested at the University of Munich Observatory, using the optical
lab set up by the OmegaCAM consortium for filter testing. A photo of the
filter at that time is shown in figure 7. The transmission of each filter
segment was measured at 21 positions forming a coarse radial pattern (fig 8)
using a monochromator beam adjusted to emulate the f-ratio 5.5 VST/OmegaCAM
optical system. The logic of the chosen measurement pattern is to give a good
sampling of the dominantly radial variation ofthe transmission profile due to
the turntable rotation in the filter coating chamber. The diameter of the
monochromator beam used in the measurements was 4-5 mm. This is a more compact
beam than that of starlight at the telescope, which fills a spotsize of up to
12 mm on passing through the filter out of focus. Consequently, the actual
performance will be a somewhat areally-smoothed version of the performance
revealed by the lab measurements and their subsequent simulation. The filter
was shipped to Paranal and VST in the spring of 2011, after some final
selective remeasuring. These confirmed there had been no discernible bandpass
changes in store since delivery almost 2 years earlier.

Figure 8: A map of the positions within each filter segment at which the
transmission was measured. The colour codings are used again in later figures
to distinguish the central positions (black, crosses, sampling roughly 30% of
the segment area), intermediate radii (blue, encircled crosses, sampling
$\sim$half the area), and corners (red crosses, 15% of the area – of which
almost half is lost to vignetting). The dashed lines define the limits of the
strips 2 arcmin wide that experience any vignetting due to the filter T-bars.
They are drawn here as for segment D.

At the time the monochromator measurements were made, a segment naming scheme
was put in place (segments A, B, C and D) which is re-used here. Presently the
filter is housed in magazine B of OmegaCAM, which means that in terms of the
view of the sky, segment A spans the SW section of the image plane, B the SE,
while C and D span the NE and NW respectively. Table 2 identifies the mosaic
CCDs beneath each segment, and sets down the centre-to-corner range in central
wavelength (CWL) and the typical throughput integral. The laboratory tests
showed us that the CWL of segments A, C and D is shortest in the segment
centre, and drifts longwards according to a centro-symmetric pattern, as the
corners and sides are approached. For segment B, the centre-to-corner drift is
reversed, with the result that the corner CWLs are bluer than in the centre of
the glass. Segment B also has the highest mean FWHM, and highest average peak
transmission: integrated over the bandpass this is a difference in throughput
of 0.045 magnitudes relative to A, C and D. The pipeline-applied illumination
correction aims to eliminate this contrast. Area-weighted transmission
profiles for the 4 segments are shown in fig 9, along with the overall mean
profile. The latter is also given numerically in table 3.

Figure 9: The mean transmission profiles of the individual glass segments, A to D (cyan, blue, green and red respectively), making up the $H\alpha$ filter and the overall mean profile (black). Table 3: Mean transmission for NB-659 Wavelength | Transmission | Wavelength | Transmission
---|---|---|---
(Å) |  | (Å) |
6456.3 | 0.000 | 6591.5 | 0.962
6461.3 | 0.001 | 6696.5 | 0.961
6465.8 | 0.001 | 6601.0 | 0.960
6470.8 | 0.002 | 6616.0 | 0.955
6475.9 | 0.002 | 6611.0 | 0.945
6480.9 | 0.003 | 6616.1 | 0.928
6485.9 | 0.005 | 6621.1 | 0.896
6490.9 | 0.008 | 6626.1 | 0.839
6496.0 | 0.012 | 6631.1 | 0.736
6501.0 | 0.020 | 6636.2 | 0.609
6506.0 | 0.033 | 6641.2 | 0.466
6511.1 | 0.053 | 6646.2 | 0.230
6516.1 | 0.086 | 6651.2 | 0.219
6521.1 | 0.136 | 6656.3 | 0.133
6526.1 | 0.208 | 6661.3 | 0.081
6531.2 | 0.307 | 6666.3 | 0.048
6536.2 | 0.429 | 6671.3 | 0.029
6541.2 | 0.575 | 6676.4 | 0.018
6546.3 | 0.700 | 6681.4 | 0.011
6551.3 | 0.799 | 6686.4 | 0.007
6556.3 | 0.868 | 6691.5 | 0.005
6561.4 | 0.915 | 6696.5 | 0.003
6566.4 | 0.936 | 6701.5 | 0.002
6571.4 | 0.947 | 6706.5 | 0.002
6576.4 | 0.954 | 6711.5 | 0.001
6581.5 | 0.959 | 6716.5 | 0.001
6586.5 | 0.961 | 6721.6 | 0.000

Compared to the H$\alpha$ filter used in the IPHAS survey, NB-659 has a CWL
that is redder on average by $\sim 20$ Å, it is around 10 percent wider, and
has a higher overall throughput leading to zeropoints $\sim$0.2 higher. The
known variations of bandpass across the 4 segments has implications for how
best to exploit VPHAS+ data. To anticipate these we have carried out two types
of simulation based on the lab measurements in order to identify them. We
describe these next, and summarise the implications in Section 3.4.

###  3.2 Simulation of the main stellar locus in the $(r-H\alpha,r-i)$ diagram

To gain an impression of the extent of the uniformity of performance with
regard to normal main sequence stars, ($r-H\alpha,r-i$) tracks were (1)
computed for each measured $H\alpha$ transmission profile using exactly the
same method as was followed by Drew et al (2005) for the analysis of IPHAS
data, (2) rescaled to a common integrated throughput, mimicking the effect of
the pipeline illumination correction, (3) compared to the mean pattern by
subtracting off the computed mean track. The result of this is shown as fig
10. The track differences picked out in red are from the segment corners
exhibiting the largest CWL shifts. It can be seen that the tracks follow the
same trend to within $\pm$0.02 up to about $r-i=1.2$ (corresponding to M3
spectral type), after which there is a clear fanning out. This shows that the
obtained $r-H\alpha$ excesses should fall within the target photometric
precision range of the survey for all except mid- to late-M stars.

The sensitivity of the M stars to variations in the narrowband transmission
profile is a point of note, while not actually a surprise. It arises from the
great breadth of the feature in M-star spectra created by the absorbing TiO
bands displaced to either side of the narrow H$\alpha$ bandpass, and the fact
that the resulting inter-band flux maximum falls at wavelengths shortward of
H$\alpha$. As these molecular bands strengthen with increasingly late spectral
type, the $r-H\alpha$ apparent excess grows along with the sensitivity to the
exact placement of the bandpass. Viewed in these terms, the $(r-H\alpha,r-i)$
colours of unreddened mid- to late-M dwarfs provide an empirical gauge of
filter bandpass uniformity and/or typical CWL. To minimise the bandpass
sensitivity and hence spread seen at late M, the CWL would need to be lowered
to around 6530 Å or less. In practice, later M dwarfs are sufficiently faint
that they normally appear in $(r-H\alpha,r-i)$ diagrams as a relatively sparse
distribution of scarcely reddened objects – falling within a thinly-populated,
continuation of the unreddened main sequence, redward of $(r-i)=1.0$, rising
from $r-H\alpha\sim 0.5$ up to $\sim 0.8$, (see figs 17, or 20). Reddened M
dwarfs are usually just too faint to be detected.

Figure 10: The r - H$\alpha$ deviations computed for all measured positions on
the H$\alpha$ filter, NB-659, after correcting the data for bandpass
integrated throughput variations. The data are colour-coded according to
position of measurement as in fig 8. The most discrepant corner positions, all
plotted in red, are located in segments C and D.

Selection of mid-to-late M dwarfs is therefore straightforward, but
quantitative interpretation of $(r-H\alpha)$ should be presumed more uncertain
than at earlier spectral types. Similar effects will be seen in the M-giant
spur located at lower $(r-H\alpha)$ in the $(r-H\alpha,r-i)$ diagram (see fig
16). However, as red giants will be picked up by VPHAS+ at large distances
through significant reddening, a precautionary check on the impact of non-zero
extinction on this fanning in colour has been made: tracks of the type
compared in fig 10 were recalculated for $A_{V}=6$ and no noticeable
additional effect was found (see also fig 17).

###  3.3 Simulation of the impact of source radial velocity on in-band
emission line fluxes

Simulations have also been performed to consider how the filter captures
emitted $H\alpha$ flux, as a function of location within the field of view and
source radial velocity. An ideal filter, centred on the mean rest wavelength
of the imaged $H\alpha$ emission and placed in a high f-ratio optical path,
would be insensitive to radial velocity shifts up to a limit proportional to
the FWHM of the bandpass. The desired capabilities of the VPHAS+ H$\alpha$
filter are separation between H$\alpha$ emission line objects and the main
stellar locus – and, better still, a regular mapping of measured $r-H\alpha$
excess onto emission equivalent width (cf Drew et al 2005, figure 6). The two
representative spectra used to investigate how these capabilities are affected
by changing source radial velocity are shown in fig 11.

Figure 11: Top panel: an example of a very bright, simple H$\alpha$ emission
profile (taken from Corradi et al 2010). The emission equivalent width is 220
Å, and the FWHM of the observed profile is close to 390 km s-1. The mean
radial velocity of the line is $+$35 km s-1. The difference between a pure
continuum magnitude and that including the line is somewhat in excess of 1.2.
Lower panel: the contrast of the line relative to continuum is much less here
(EW $\sim$ 20 Å), and the FWHM is somewhat wider at 570 km s-1 (a classical Be
star, taken from Raddi et al, 2013). The mean radial velocity is $-$50 km s-1.
The continuum-only magnitude is fainter by about 0.2 only, here.

Both spectra were blueshifted to $-$500 km s-1 , and then shifted redward in
steps of 100 km s-1 at a time, up to $+$500 km s-1 (altogether a displacement
of 22 Å) – calculating at each step the integral of the spectrum folded
through the filter transmission profile. The resultant in-band fluxes were
converted to magnitudes, and then shifted by the amount required to match the
integrated transmission to the overall mean for the filter (again mimicking
the function of the pipeline illumination correction). In real use, we would
expect the majority of emission line objects to present with FWHM no greater
than either of these examples (interacting binaries and WR stars do present
with much broader emission, however).

The radial velocity range explored was chosen with the following
considerations in mind:- Emission line stars in the thin disk will commonly
have radial velocities falling within the range -100 to +100 km/s. In the
Bulge larger radial velocities may be encountered: excursions to $\pm$200 km
s-1 are observed in CO (Dame, Hartmann & Thaddeus 2001) within $\sim
20^{\circ}$ of longitude of the Galactic Centre, and for a minority of inner-
Galaxy planetary nebulae, radial velocities have been obtained that extend the
range almost to $\pm$300 km s-1 (see Durand, Acker & Zilstra 1998, and
Beaulieu et al 2000).

Figure 12: Results of simulation of the in-band flux as a function of source
radial velocity for the EW = 220 Å emission spectrum (upper panel in fig 11).
The fluxes are expressed as magnitude offsets relative to the peak simulated
in-band flux. The panels representing the segments are arranged as they are
imposed on the plane of the sky, i.e. A covers the SW quadrant, D covers the
NW when the filter is stored in OmegaCAM’s magazine B. The curves are colour-
coded according to measurement position as in fig 8. Galactic sources will
usually fall well within the marked velocity range, $-$200 to $+$200 km s-1.
The more problematic red curves, representing the response of the segment
corners, account for $\sim$8% of each segment’s unvignetted area only. Figure
13: Results for an emission line net equivalent width of 20 Angstroms
(spectrum shown in lower panel in fig 11). Otherwise as fig 12. Figure 14: The
layout of the H$\alpha$ filter, showing the positions of the PN-test
measurements, colour coded according to flux relative to the corners of
segment B. Those shown as black crossed boxes have a scaled flux $\geq 0.95$.
Those in blue are a little lower with a scaled flux in the range 0.90 – 0.94.
The two positions marked in green have fluxes scaling to 0.88 and 0.89, while
the two in red scale to 0.84 and 0.77 (both in segment D). There is no
measurement for the centre of segment C because of a telescope pointing
imprecision. The other centre measurements were obtained $\sim 2$ arcmins away
from the true centres plotted in order to avoid CCD gaps. The dashed lines
mark the limits of the T-bar vignetting.

The results of this exercise are plotted in figs 12 and 13. Fig 12 shows that
segments A and B come very much closer to independence of radial velocity than
segments C and D, in terms of measured $H\alpha$ excess. In all segments,
reasonable fidelity (a flat, or nearly flat response) is achieved around
segment centre – although in B, uniquely, the corners happen to perform a
little better than the centre. Clearly segment D, where the transmission is
centred on longer wavelengths than in the other segments, could yield
measurements in its corners (perhaps 8% of its unvignetted area) of $H\alpha$
magnitude, or of $r-H\alpha$, that underestimate the flux of similarly high
equivalent-width H$\alpha$ emission by up to $\sim$ 0.3 (out of a true excess,
expressed in magnitudes of $\sim 1.2$). Segment C performs similarly, but the
potential flux drop associated with its corners is less pronounced.

We can compare the expectations created by these simulations with the results
of an on-sky experiment in which the planetary nebula (PN) ESO 178-5 (or PNG
327.1 -02.2) has been exposed in H$\alpha$ at a series of positions in the
image plane, placing it well into every corner and also close to the centre of
each of the four filter segments (see fig 14). The observed integrated counts
variation might be predicted to be somewhat stronger than in fig 12 given that
the in-band continuum flux from this PN will be relatively even weaker. But
this will be offset by the additional flux due, in particular, to [NII]
$\lambda$6584\. The PN chosen for this test was picked both because it is
well-calibrated (Dopita & Hua 1997), and because its LSR radial velocity is
quite large and negative ($v_{LSR}=-88.7$ km s-1: on 19th April 2013 when it
was observed, this will have shifted to $-105$ km s-1 at the telescope). It
also happens to possess [NII] $\lambda$6584 emission that is scarcely less
bright than H$\alpha$ (the former has 98% the flux of the latter: Dopita & Hua
also provide a spectrum of this nebula).

Background-subtracted aperture photometry of the PN and a moderately bright
star nearby, serving as a continuum reference, was carried out on the reduced
images. These measurements reveal a pattern of behaviour that essentially
tracks the results shown in figure 12: the continuum reference itself shows a
total count variation of $\pm 5$% across all pointings, while the PN counts,
after scaling to the reference, range from $+5$%, down to $-23$% relative to
the values for the corners of segment B. In the extreme case that all of the
H$\alpha$ and [NII] $\lambda$6548 emission had been shifted out of the
bandpass, the maximum drop for this PN would be $-57$% (the remaining 43%
being attributable to [NII] $\lambda$6584). The pattern across the filter of
the results emerging from this trial is shown in fig 14.

The radial-velocity dependence of the transmitted flux may accordingly become
an issue for objects with very strong H$\alpha$ emission where the aim is
accurate flux determination, unless attention is paid to where the object
falls in the image plane. Qualitatively the issue is less critical: regardless
of where the object is located, the changes in transmission are not so large
that there will be frequent failures to distinguish strong H$\alpha$ emitters
– i.e. they will still appear above the main stellar locus in the
$(r-H\alpha,r-i)$ diagram. In the example simulated, the outer reaches of
segment D would bring $(r-H\alpha)$ down to 0.9 – 1.0, a level that
nevertheless remains clear of the domain that might be occupied by unreddened,
non-emission very late-type M dwarfs ($(r-H\alpha)\sim 0.8)$, cf fig 16). As
further context, we note that nearly continuum-free emission line objects,
such as PNe and HII regions, present with $r-H\alpha\sim 3$.

Where the line emission itself contributes only a minority of the measured
narrowband $H\alpha$ flux the trends seen are much more subdued (fig 13). For
the example shown, only 20 percent of the total in-band flux is attributable
to the net line emission, rather than most of it as in fig 12. Again, the
corners of segments C and D perform least well, in under-representing the
emission flux by up to $\sim$0.05 magnitudes at the most negative likely
Galactic Plane radial velocities. Otherwise, the performance is predicted to
be within the anticipated 0.02–0.03 error budget of the survey.

###  3.4 Implications of the H$\alpha$ filter properties for VPHAS+ and its
exploitation

In summary, for most purposes the H$\alpha$ filter performs as required, and
has very good throughput. For the great majority of stars making up the main
stellar locus, there will be the desired fidelity of $(r-H\alpha)$ colour, and
the great majority of emission line objects will be detected with the same
facility as they are by IPHAS.

There are two caveats to note. First, in Section 3.2 it was shown that
variations in central wavelength across the filter segments will lead to
thickening of the loci traced by mid-to-late M stars. These same variations,
of what is a relatively red H$\alpha$ passband, also introduce the potential
for under-determination of H$\alpha$ fluxes for objects/nebulosity in parts of
the image plane for sources with significantly negative radial velocities
(Section 3.3). This becomes most serious for emission-line sources falling
near the vignetted corners of segments D and C, where a $20-30$ percent under-
counting in H$\alpha$ may occur for radial velocities approaching $-200$ km
s-1. As is always the case for narrowband H$\alpha$ filters, the common
presence of significant [NII] $\lambda\lambda$6548, 6584 emission bracketing
H$\alpha$ in planetary nebulae or HII regions complicates the expected
signature. However, it can be guaranteed in all but rare, exotic circumstances
that the stronger $\lambda$6584 component of the [NII] doublet falls well
within the bandpass.

The outstanding practical consequence of the filter’s transmission
characteristics for the survey strategy are that, for quantitative
reliability, measures of $(r-H\alpha)$ obtained using segments A and B, and
the central zones of C and D (out to $\sim 13$ arcmin) are to be favoured.
This appreciation is half of the reason for the adoption of offsets of several
arcminutes between the 3 successive pointings made in this filter (fig 3) for
each field – our strategy ensures that objects captured in segment corners in
the first pointing are fall close to segment centres in the third.

The rest of the motive is to mitigate the cross-shaped vignetting due to the
blackened T-bars holding the segments in place (fig 7). Each arm of the cross
casts a shadow entirely contained within a strip 4 arcmin wide. By choosing to
offset at least this much in both RA and Dec between the 3 exposures obtained
per field, we raise the probability of at least one high-confidence
$r-H\alpha$ colour measurement per detected source in the final catalogue to
very nearly 100 percent, and the probability of two to over 95 percent.

Finally, we remark that the combination of large offsets and three pointings
has the consequence that the fraction of sky within the survey footprint
missed altogether, due to dead areas between the CCDs and vignetting, is under
0.3 percent.

##  4 Simulation of VPHAS$+$ stellar colours

The five bandpasses of the survey provide the basis for the construction of a
range of magnitude-colour and colour-colour diagrams. To take full advantage
of them, knowledge is needed of the behaviours that can be expected of the
colours of normal stars.

We have simulated colours for solar-metallicity main sequence and giant stars
using the same method as employed by Sale et al. (2009). We adopt the
definition of these two sequences in $(T_{eff},\log g)$ space given by
Straizys & Kuriliene (1981). Then for each spectral type along these
sequences, solar-metallicity model spectra were drawn from the Munari et al.
(2005) library. At a binning of 1 Å the spectra in this library are well
enough sampled to permit the calculation of narrow-band $H\alpha$ relative
magnitudes with confidence, alongside the analogous broadband quantities. More
detail on the broadband filter transmission profiles, shown in Fig. 2, and on
the CCD response curve is provided on the ESO
website444http://www.eso.org/sci/facilities/paranal/instruments/omegacam/doc.
To ensure compliance with the Vega-based zero magnitude scale, we have defined
the synthetic colour arising from a flux distribution $F_{\lambda}$ as
follows:

$m_{1}-m_{2}=-2.5\log\left[\frac{\int T_{1}\lambda F_{\lambda}d\lambda}{\int
T_{1}\lambda F_{\lambda,V}d\lambda}\right]+2.5\log\left[\frac{\int
T_{2}\lambda F_{\lambda}d\lambda}{\int T_{2}\lambda
F_{\lambda,V}d\lambda}\right]$

where $T_{1}$ and $T_{2}$ are the numerical transmission profiles for filters
$1$ and $2$, after multiplying them through by the atmospheric transmission
(Patat et al 2011) and mean OmegaCAM CCD response curves. The SED adopted for
Vega, $F_{\lambda,V}$, is that due to Kurucz
(http://kurucz.harvard.edu/stars.html). Where needed for comparison, we have
also computed colours based on the Pickles (1998, hereafter P98)
spectrophometric stellar library (the approach adopted by Drew et al 2005 for
IPHAS). To maintain precision, the numerical quadrature resamples the more
smoothly varying transmission data onto the sampling interval of the stellar
SED.

The $(r-H\alpha)$ excess is evaluated in exactly the same way as the broadband
colours. Since Vega is an A0V star, its SED at H$\alpha$ incorporates a strong
absorption line feature that reduces the in-band flux below the pure continuum
value. Unlike the broadbands, the $H\alpha$ narrowband has not yet been
standardised and so there is not a formally recognised flux scale. However, we
can specify here that the integrated in-band energy flux for Vega, on adopting
the mean profile for the VST filter, is $1.84\times 10^{-7}$ ergs cm-2 s-1 (at
the top of the Earth’s atmosphere). To assure zero colour relative to the
optical broad bands, this flux is required to correspond to $m_{H\alpha}\simeq
0.03$. The reduction in zeropoint (zpt) that the computed in-band flux implies
relative to the flux captured by the much broader $r$ band – based on folding
Vega’s SED with lab measurements of the filter throughputs corrected for
atmosphere and detector quantum efficiency – is 3.01. Current practice in
VPHAS+ photometric calibration is accordingly to adopt zpt(NB-659) = zpt(r) -
3.01 magnitudes as the default calibration for the narrowband: in section 5
where a direct comparison is made with SDSS spectroscopy, this offset is found
to be satisfactory. When applied, it assures that data obtained in
photometric, or stable, conditions, yield zero $r-H\alpha$ colour for A0
stars.

Figure 15: The expected positions of main sequence and giant stars in the
$(u-g,g-r)$ plane. For the main sequence, tracks are shown for the
monochromatic reddenings $A_{0}=$ 0, 2, 4, 6 and 8 (working from left to
right). The red leak in the $u$ filter starts to lower $u-g$, noticeably from
$A_{0}=6$ (tracks drawn in red). The giant-star tracks, drawn as dashed lines
for $A_{0}=0$, 2 and 4 only, are very similar to their main-squence
counterparts except at the latest types. Figure 16: The expected position of
main sequence and giant stars in the $(r-H\alpha,r-i)$ plane. Tracks are shown
for the monochromatic reddenings $A_{0}=$ 0, 2, 4, 6, 8 and 10 (from left to
right). Solid lines represent the main sequence tracks, while dashed lines are
used for the giant tracks. The lines in red are giant-star tracks derived from
P98 spectrophotometry.

Both main-sequence and giant-star colours have been calculated for a range of
reddenings and optical-IR extinction laws as formulated by Fitzpatrick & Massa
(2007). The unreddened colours for the mean Galactic law ($R=3.1$) are laid
out here in table 4. The Appendix provides additional tables that specify the
colours of main sequence stars at selected reddenings and for two further
representative reddening laws ($R=2.5$ and 3.8). These can be used to
construct intrinsic-colour-specific reddening lines. For the large range in
extinction sampled along many Galactic Plane sightlines, these reddening
trends are slightly curved (see the examples shown in e.g. Sale et al 2009).
In this paper we use $A_{0}$, the monochromatic reddening at 5500 Å to
parameterise the amount of reddening, rather than the band-averaged measure,
$A_{V}$. In most circumstances these quantities are almost identical.

Table 4: VST/OmegaCAM synthetic colours for unreddened main-sequence dwarfs and giants. Sp.type | main sequence (V) |  | giants (III)
---|---|---|---
|  |  |  |  |  | (model) | (P98 library spectra)
| $u-g$ | $g-r$ | $r-i$ | $r-H\alpha$ |  | $u-g$ | $g-r$ | $r-i$ | $r-H\alpha$ | $r-i$ | $r-H\alpha$
O6 | -1.494 | -0.313 | -0.145 | 0.071 |  |  |  |  |  |  |
O8 | -1.463 | -0.299 | -0.152 | 0.055 |  |  |  |  |  | -0.158 | 0.074
O9 | -1.426 | -0.271 | -0.142 | 0.064 |  | -1.426 | -0.271 | -0.142 | 0.064 |  |
B0 | -1.404 | -0.267 | -0.143 | 0.058 |  | -1.404 | -0.267 | -0.143 | 0.058 |  |
B1 | -1.296 | -0.236 | -0.130 | 0.052 |  | -1.316 | -0.234 | -0.130 | 0.057 | -0.095 | 0.071
B2 | -1.181 | -0.214 | -0.117 | 0.049 |  | -1.209 | -0.211 | -0.116 | 0.056 |  |
B3 | -1.025 | -0.182 | -0.098 | 0.048 |  | -1.046 | -0.182 | -0.098 | 0.054 | -0.035 | 0.083
B5 | -0.799 | -0.133 | -0.071 | 0.043 |  | -0.814 | -0.134 | -0.072 | 0.050 | -0.016 | 0.083
B6 | -0.699 | -0.116 | -0.062 | 0.040 |  | -0.714 | -0.116 | -0.062 | 0.046 |  |
B7 | -0.550 | -0.094 | -0.051 | 0.033 |  | -0.568 | -0.095 | -0.051 | 0.041 |  |
B8 | -0.361 | -0.071 | -0.039 | 0.022 |  | -0.383 | -0.072 | -0.039 | 0.032 |  |
B9 | -0.168 | -0.040 | -0.023 | 0.009 |  | -0.186 | -0.044 | -0.024 | 0.021 | -0.018 | 0.035
A0 | -0.024 | 0.000 | -0.003 | -0.002 |  | -0.030 | -0.009 | -0.006 | 0.011 | 0.012 | 0.034
A1 | 0.007 | 0.015 | 0.004 | -0.004 |  | 0.007 | 0.004 | 0.000 | 0.008 |  |
A2 | 0.039 | 0.038 | 0.014 | -0.005 |  | 0.051 | 0.022 | 0.009 | 0.008 |  |
A3 | 0.064 | 0.062 | 0.025 | -0.005 |  | 0.085 | 0.045 | 0.019 | 0.007 | 0.037 | 0.063
A5 | 0.096 | 0.130 | 0.056 | 0.008 |  | 0.143 | 0.107 | 0.048 | 0.015 | 0.096 | 0.087
A7 | 0.073 | 0.206 | 0.089 | 0.030 |  | 0.145 | 0.179 | 0.078 | 0.032 | 0.115 | 0.094
F0 | 0.003 | 0.336 | 0.153 | 0.086 |  | 0.091 | 0.317 | 0.144 | 0.084 | 0.156 | 0.102
F2 | -0.021 | 0.396 | 0.182 | 0.111 |  | 0.064 | 0.380 | 0.174 | 0.109 | 0.204 | 0.172
F5 | -0.039 | 0.505 | 0.230 | 0.150 |  | 0.046 | 0.491 | 0.224 | 0.148 | 0.238 | 0.164
F8 | -0.013 | 0.587 | 0.263 | 0.174 |  |  |  |  |  |  |
G0 | 0.012 | 0.628 | 0.278 | 0.185 |  |  |  |  |  | 0.333 | 0.213
G2 | 0.011 | 0.628 | 0.279 | 0.185 |  | 0.253 | 0.759 | 0.329 | 0.215 |  |
G5 | 0.162 | 0.756 | 0.327 | 0.217 |  | 0.405 | 0.870 | 0.368 | 0.235 | 0.396 | 0.250
G8 | 0.355 | 0.845 | 0.358 | 0.233 |  | 0.531 | 0.944 | 0.395 | 0.247 | 0.419 | 0.247
K0 | 0.523 | 0.938 | 0.396 | 0.248 |  | 0.640 | 1.002 | 0.417 | 0.256 | 0.446 | 0.254
K1 | 0.551 | 0.954 | 0.403 | 0.251 |  | 0.803 | 1.085 | 0.451 | 0.269 | 0.468 | 0.269
K2 | 0.629 | 0.993 | 0.419 | 0.258 |  | 0.963 | 1.159 | 0.483 | 0.281 | 0.508 | 0.293
K3 | 0.779 | 1.062 | 0.447 | 0.269 |  | 1.227 | 1.276 | 0.536 | 0.299 | 0.514 | 0.286
K4 | 0.871 | 1.108 | 0.468 | 0.278 |  | 1.374 | 1.342 | 0.585 | 0.320 | 0.592 | 0.313
K5 | 1.083 | 1.210 | 0.522 | 0.300 |  | 1.578 | 1.420 | 0.630 | 0.338 | 0.714 | 0.337
K7 | 1.387 | 1.402 | 0.724 | 0.387 |  |  |  |  |  |  |
M0 | 1.372 | 1.411 | 0.789 | 0.411 |  | 1.697 | 1.454 | 0.686 | 0.360 | 0.827 | 0.411
M1 | 1.335 | 1.439 | 0.934 | 0.467 |  | 1.838 | 1.506 | 0.769 | 0.395 | 0.872 | 0.401
M2 | 1.262 | 1.442 | 1.112 | 0.522 |  | 1.938 | 1.546 | 0.825 | 0.420 | 0.920 | 0.443
M3 | 1.236 | 1.447 | 1.179 | 0.545 |  | 1.980 | 1.556 | 0.933 | 0.444 | 1.165 | 0.471
M4 | 1.248 | 1.457 | 1.168 | 0.543 |  | 1.959 | 1.551 | 1.136 | 0.500 | 1.472 | 0.512
M5 |  |  |  |  |  | 2.009 | 1.569 | 1.296 | 0.556 | 1.739 | 0.560
M6 |  |  |  |  |  | 2.199 | 1.612 | 1.267 | 0.554 |  |

Based on the data from these tables, the main sequence and giant tracks are as
shown in figures 15 and 16. These identify where the main stellar loci will
fall. It is important to note that the OmegaCAM $u$ filter, like all filters
constructed for this challenging band, exhibits a low-level red leak. In this
instance, lab measurements indicate transmission at levels between $10^{-5}$
and $10^{-4}$ within limited windows around $\sim$9000 Å. This is enough to
begin to noticeably, and erroneously, brighten the $u$ magnitudes of normal
stars reddened to $g-r>3$. Because of this, and because the measurement of
very low level leakage is itself subject to proportionately higher
uncertainty, we do not plot or tabulate $u-g$ data beyond $g-r=3$ limit. Very
few detected sources are so extreme. In practice, VPHAS+ $u-g$ is faithful for
extinctions up to $A_{0}\sim 6$, but gradually thereafter it transforms into a
colour that behaves crudely as $-(g-z)$.

The $(r-H\alpha,r-i)$ colour-colour diagram is not subject to such effects,
and therefore remains sound across a wider spread in visual extinctions.
Synthetic tracks are presented in figure 16 for $A_{0}=0$, 2, 4, 6, 8 and 10.
The main sequence tracks shown are similar to those appropriate to IPHAS (cf
figure 6 of Sale et al 2009). But a problem emerges when it comes to the
simulation of red giant colours. Purely theoretical simulation predicts late-K
and M giant colours closely resembling those of dwarfs, whereas simulation
using P98 library spectrophotometry indicates a distinctive flattening of the
M-giant track, peeling away from the steadily rising main sequence track.
Figure 16 points out this contrast. Inspection of table 4 reveals this is a
problem linked mainly to simulation of the $i$ spectral range, which renders
$(r-i)$ progressively larger from late K into the M giant range when the
library spectra are used in place of model atmospheres.

Figure 17: Equatorial VPHAS+ data (upper panel) and IPHAS data (lower panel)
compared to show the different appearance of M-giant $(r-H\alpha,r-i)$
colours. The photometry is extracted from a $\sim$0.2 sq.deg region of sky,
centred on $\ell=35.95^{\circ}$, $b=-3.13^{\circ}$. The magnitude range is
limited to $13<r<18$ in both cases. Telescope-appropriate giant tracks,
computed from the P98 spectrophotometric library, for $A_{0}=$ 2 and 4 are
superimposed in red in both panels. In the upper panel, giant tracks computed
from model atmospheres for the same reddenings are shown in blue. The black
dashed line in each panel is the synthetic early-A reddening line, from
$A_{0}=0$ to $A_{0}=10$. The stars appearing in the IPHAS catalogue as
candidate emission line stars are enclosed in cyan boxes in both panels.

Evidence that M giants are better reproduced by synthetic photometry based on
flux-calibrated spectra is provided by figure 17. This figure also compares
new VPHAS+ data with their crossmatches in the IPHAS survey within a $\sim$0.2
sq.deg equatorial field ($\ell=35.95^{\circ}$, $b=-3.13^{\circ}$), and shows
selected synthetic tracks superimposed. The photometry from the two surveys of
the most densely populated part of the main stellar locus to $(r-i)\simeq 1.5$
substantially overlap, but not perfectly – the response functions describing
the three bandpasses involved in fig 17 undoubtedly differ in detail between
the two telescopes. On cross-correlating either $(r-i)$ or $(r-H\alpha)$
between the two surveys, it becomes clear that the IPHAS colour has the
somewhat larger dynamic range. This is the reason for the slightly more
stretched appearance of both the main locus and the early-A reddening line in
the IPHAS diagram relative to that for VPHAS+.

At $(r-i)>1.5$ in fig 17, it can be seen that the M-giant spurs look very
different. First, the VPHAS+ M giants fall into a nearly flat distribution
lying at lower $(r-H\alpha)$, compared to the more steeply rising higher IPHAS
M-giant sequence. However, as long as the data are interpreted with reference
to telescope-appropriate synthetic photometry, the two datasets will lead to
the same inference. In the example shown in fig 17, the comparisons with
suitable synthetic giant tracks indicate that the maximal extinction in the
field can be no more than $A_{0}\simeq 4$. The extinction measures due to
Marshall et al (2006), based on 2MASS red-giant photometry, indicate a maximum
Galactic extinction of $A_{K}\sim 0.3$ for this pointing. For a typical
Galactic $R=3.1$ reddening law this scales up to $A_{0}\sim 3.3$ (roughly –
see Fitzpatrick & Massa 2009). If model-atmosphere giant tracks are referred
to instead, the M giants would have to be read as demanding visual extinctions
ranging from $\sim$4 upwards.

Fig 17 also demonstrates the broadening in $(r-H\alpha)$ of the VPHAS+ M-giant
sequence that was foretold in Section 3. The IPHAS counterpart is evidently
much sharper, as it rises to higher $r-H\alpha$ with increasing $r-i$. The
main practical impact of this difference is that IPHAS M-giant photometry is
the better starting point for picking apart chemistry differences (Wright et
al 2008). But it is as true of the VPHAS+ $(r-H\alpha,r-i)$ diagram as it is
of its IPHAS equivalent – that M giants at $r-i>1.5$ sit below and apart from
M dwarfs.

Fortuitously, fig 17 identifies an advantage of the generally good seeing
available at the VST. There are two candidate emission line objects apparent
in the IPHAS selection (enclosed in cyan boxes, in fig 17), that drop back
into the main stellar locus in the VPHAS+ data. Inspection of the images shows
that both stars are in close doubles of similar brightness, of under 2 arcsec
separation. Because they are a little better resolved in VPHAS+ ($\sim$0.8
arcsec seeing), than in IPHAS ($\sim$0.9 arcsec seeing), the pipeline makes a
better job of the assigning magnitudes in the different bands to the blend
components. This example nicely illustrates the most common reason for bogus
candidate emission line stars in either VPHAS+ or IPHAS - improperly
disentangled blends. Candidate emission line stars should always be checked
for this kind of problem before spectroscopic follow-up. Otherwise, experience
with IPHAS gives confidence that the selection of emission line objects via
VPHAS+ will be highly efficient (see e.g. Vink et al 2008, Raddi et al 2013).

Finally, it is worth noting that the bright limit of the survey at 12–13th
magnitude effectively excludes any unreddened stars of earlier spectral type
than $\sim$G0. Before more luminous stars of spectral type F and earlier can
enter the survey sensitivity range, they need to be at distances in excess of
1 kpc, typically, where low extinction becomes increasingly improbable. This
constraint bestows a significant selection benefit in that only unreddened or
lightly-reddened subluminous objects, with intrinsically blue colours are left
standing clear near the blue end of the main stellar locus in commonly-
constructed photometric diagrams. In this domain VPHAS+ has important
selection work to do.

##  5 Photometry validation: a comparison of SDSS and VST data

Before the start of survey field acquisition, we obtained observations in all
survey filters of two pointings that fall within the SDSS photometric and
spectroscopic coverage (Abazajian et al 2009). These were centred on RA
20:47:53.7 Dec -06:04:14.5 (J2000) and RA 21:04:25.94 Dec +00:59:15.8 (J2000)
– fields that happen to include a number of white dwarfs and cataclysmic
variables (not discussed further here). The main aim of the data was to verify
VST photometry both by comparison to SDSS photometry and to synthetic
photometry derived from SDSS spectra. The VST observations were obtained on
21/09/2011, during clear weather at a time of generally sub-arcsecond seeing.
The exposure times differ only a little from those now in general survey use:
the $g$ exposures were 30 sec, rather than 40, and $i$ was exposed for 20 sec
rather than 25 - the other times were as given in Section 2.1.

The photometry on the sources in these fields have been pipeline-extracted and
calibrated in the standard way, and have been cross-matched to their SDSS
counterparts. The number of cross matched stars used in this exercise ranges
from $\sim 2500$ ($u$) up to $\sim 10000$ ($r$). For a star to be included, it
must be: unvignetted; to have a star-like point spread function; to lie within
0.5 arcsec of its SDSS counterpart, and to fall within the magnitude range
$16>r_{VST}>19$. The SDSS selection constraints were set to exclude blended
and saturated sources, and sources close to detector edges. In addition, it
was required of every source that, in both surveys, the formal error on the
magnitude measurement is less than 0.03.

Figure 18: The distributions, by band pass, of the measured magnitude
differences between VST test data and SDSS, after correction of the latter
magnitudes from the AB to the Vega zero-magnitude scale. These were obtained
from two VST pointings, away from the Galactic Plane, that overlap the SDSS
footprint. The biggest differences are found in the $u$ band where the VST
data are fainter than the corrected SDSS values by $\sim 0.12$, in the median.
Figure 19: Measured magnitude shifts between VST and SDSS cross-matched
objects as a function of SDSS colour, for the second of the two fields
observed (only). The data in blue are $\Delta u$ versus $(u-g)_{SDSS}$; the
data in green and red are, respectively, $\Delta g$ and $\Delta r$ versus
$(g-r)_{SDSS}$; the data in black are $\Delta i$ versus $(r-i)_{SDSS}$. The
horizontal lines show where the four loci would be expected to lie in the case
that the SDSS and VST broadband filters were identical, and the calibrations
perfect. The data fall into loci that are not far displaced from these
horizontal lines and are almost as flat: this indicates that the colour-
dependent terms that would be needed in equations to transform between the
SDSS and VST systems are small.

In fig 18 we plot the histograms of the magnitude differences between the two
surveys, according to pass band – pooling the data from both pointings. If the
starting assumption is that the VST broadband filters are identical to the
SDSS set, the predicted magnitude for each star in each filter is the measured
SDSS magnitude, less the offset between the AB and Vega scales (essentially
the numerical difference between the magnitude of Vega in the AB system and
its value of 0.02 to 0.03, according to the alternative Vega-based convention
– see table 8 in Fukugita et al 1996). In the $g$, $r$ and $i$ bands, the
predicted and observed magnitudes are well-enough aligned, and the
interquartile spread is consistent with the way the data were selected for
random errors less than 0.03. However in $u$ there is a discrepancy that
exceeds expected error: the median difference deviates by 0.12 magnitudes and
the width of the distribution is twice that arising in the comparisons of the
other bands.

The fuller picture is presented in figure 19 which shows the broadband
magnitude differences as a function of the relevant SDSS colour for the second
of our two fields (only). In all 4 pass bands, including $u$, the colour
dependence can be seen to be very weak in that the loci traced out by the
plotted stars are – to a first approximation – flat. The discrepancy seen in
$u$ is revealed as mainly a zero-point shift, combined with scatter that
exceeds the formal errors. In highly reddened Galactic Plane fields, the
stellar colour effects may become more pronounced as extinction modifies the
effective sampling of the passbands.

Table 5: Mean magnitude offsets between VST photometry and cross-matching synthetic photometry derived from the SDSS database of spectra. The synthetic magnitude scale adopts magnitudes for Vega itself of 0.026 (see Bohlin & Gilliland 2004) Offset | Field 1 (117 stars) | Field 2 (50 stars)
---|---|---
$u_{VST}-u_{syn}$ | 0.07$\pm$0.39 | 0.11$\pm$0.17
$g_{VST}-g_{syn}$ | 0.06$\pm$0.09 | 0.06$\pm$0.04
$r_{VST}-r_{syn}$ | 0.01$\pm$0.06 | 0.03$\pm$0.03
$i_{VST}-i_{syn}$ | -0.06$\pm$0.08 | -0.05$\pm$0.03

As a separate exercise, we have used SDSS spectra to synthesise magnitudes and
colours for stars with cross-matching VST photometry. The spectral type range
present within this much smaller sample runs from B-type through to early
M-type (M1). At wavelengths below 3800 Å falling within the $u$ band, it was
necessary to extrapolate the spectra using appropriately chosen P98 library
data. The result of this comparison is agreement between the VST and
synthesised magnitudes at the $\sim$5 percent level (table 5), with the $u$
band as the outlier exhibiting much more pronounced scatter as well as
somewhat higher offset. This pattern echoes the behaviour apparent in the VST-
SDSS purely photometric comparison of fig 18, using a much larger sample. The
difficulty is not confined to VST $u$ however, in that SDSS $u$ photometry
fares scarcely any better relative to synthesis from the spectra (for the two
fields, offset and scatter are -0.09 $\pm$ 0.37, and -0.04 $\pm$ 0.16). As
more blue survey data are accumulated, it may become clear that the $u$
zeropoint will benefit from being tied to that of $g$ for those fields
observed in the best conditions, as is presently done for $H\alpha$ with
respect to $r$. This option is not yet enacted. For the time being, it must be
acknowledged that pipeline $u$ calibration is more approximate than those of
the other bands.

We have also used the reduced cross-match sample to look at how the VST
photometric $r-H\alpha$ colour compares with its counterpart synthesised from
spectra – looking, in particular, for any trends as a function of distance
from field centre. No such trend is apparent, thereby meeting the expectation
that the narrowband fluxes of normal stars, to early-M spectral type, would
not be affected by the pattern of bandpass shifts discussed earlier in section
3 (cf. fig 10). However we do find that in order to make this detailed
comparison, systematic offsets had to be removed from the VST photometry
first. These were 0.075 in $(r-i)$, in the sense that the VST colours were too
red by this amount, and a 0.02 reduction in $(r-H\alpha)$. The $(r-i)$ offset
is consistent with the broadband magnitude offsets listed in Table 5 and hence
is as expected. The $(r-H\alpha)$ adjustment is small enough (i.e. within the
fit error) that it supports the zeropoint shift of 3.01 magnitudes between the
$r$ and $H\alpha$ bands that was identified in section 4). Once these colour
offsets are applied to the VST data, the rms scatter of the photometric
$(r-H\alpha)$ colour relative to its synthetic counterpart is 0.04 for objects
brighter than $r=19$.

Figure 20: Left, the $(u-g,g-r)$ and right, the $(r-H\alpha,r-i)$ photometric
diagrams pertaining to VPHAS+ survey field 1679. Both diagrams are plotted as
two-dimensional stellar-density histograms, rainbow colour-coded such that
high source densities (80-90 per bin) are red and the lowest densities (one
per bin) are dark blue. The binning is 0.017$\times$0.025 in the left panel,
0.013$\times$0.008 in the right. The synthetic unreddened main sequence is
drawn in, in black, in both panels. The G0V and A3V reddening lines obtained
for $R=3.1$, drawn as black dashed lines, are included in respectively the
$(u-g,g-r)$ and $(r-H\alpha,r-i)$ diagrams as useful aids to interpretation.

##  6 An example of point-source photometry derived from a VPHAS+ field

The extracted point-source photometry from the VST square-degree field is one
of the two main data products from VPHAS+ – the other being the images
themselves, considered below in section 7. We present an example of the two
essential colour-colour diagrams in fig 20, in which band-merged stellar
photometry for field 1679 is compared with the primary diagnostic synthetic
tracks presented in section 4. This field includes the sky area from which
data were taken to construct fig 6 illustrating typical errors. The massive
open cluster, Westerlund 2, is located in the NE of these pointings, and the
field as a whole includes moderate levels of diffuse, complex Hii emission.
The data presented are drawn from a sky area centred on RA 10 24 49 Dec -57 58
00 (J2000) that spans 1.3 sq.deg – the total footprint occupied by the two
offset positions.

Only objects with stellar point-spread functions in $g$, $r$ and $i$, brighter
than $g=20$ are included in fig 20. Where two sets of magnitudes are
available, the mean values have been computed and used. A further requirement
imposed is that the random error in all bands may not exceed 0.1. The same
$\sim$37000 objects are included in both diagrams. In order to obtain the
diagrams shown, the pipeline photometric calibration was checked and refined
as follows: we

  * •

cross-matched brighter stars to APASS $g$, $r$ and $i$ photometry

  * •

computed the median magnitude offset (applying no colour corrections – it was
shown in fig 19 these are modest)

  * •

corrected all $g$, $r$ and $i$ for these offsets;

  * •

corrected $u$ by determining the vertical shift needed in the $(u-g,g-r)$
diagram to align the main stellar locus with the unreddened main sequence and
the G0V reddening line

  * •

corrected the $H\alpha$ zeropoint and hence all $H\alpha$ magnitudes according
to the requirement that $zpt(H\alpha)=zpt(r)-3.01$.

This resulted in the following broadband corrections:- $\Delta i=-0.004$,
$\Delta r=-0.032$ (red filter set), $\Delta r=-0.033$ (blue filter set),
$\Delta g=0.069$ and $\Delta u=-0.31$. As expected, the correction that had to
be applied to the $u$ photometry was, by far, the largest.

The main stellar locus can be seen to be tightly concentrated in both the blue
and the red diagrams, and to favour lightly reddened G and K stars. The
superimposed synthetic reddening lines (G0V in the $(u-g,g-r)$ diagram, A3V in
$(r-H\alpha,r-i)$) have been drawn adopting the $R=3.1$ reddening law widely
regarded as the Galactic norm. The blue diagram provides examples of three
distinct typical populations falling outside the main stellar locus. Below it,
at $(u-g)>1.5$ and $(g-r)>1.5$ (roughly) the plotted objects will mainly be M
giants. Above the main stellar locus toward the red end, in the ranges
$0<(u-g)<0.5$ and $1.5<(g-r)<2.0$ lie the OB stars in and around Westerlund 2.
Finally, the modest scatter of blue objects lying above the G0V line roughly
in the $0\leq(g-r)\leq 1$ range will include intrinsically-blue lightly-
reddened subluminous objects.

It is interesting to note in the red diagram that there is some evidence that
early-A stars making up the lower edge of the main stellar locus would better
follow a different law, with $R\sim 3.8$ (see the tables in the Appendix).
Indeed a reddening law of this type has been inferred for the OB stars in
Westerlund 2 by Vargas Alvarez et al (2013). Most of the thin scatter of
points below the main stellar locus, and some of the scatter above, in this
same diagram will be the product of inaccurate background subtraction in
H$\alpha$. But many of the objects lying above the main stellar locus will
indeed be emission line objects, and some of the stars below will be white
dwarfs. As expected, the red spurs of M dwarfs and M giants are broader
features than their IPHAS counterparts (cf fig 17 and associated remarks).

For more discussion of these colour-colour diagrams, the reader is referred to
Groot et al (2009, UVEX) and Drew et al (2005, IPHAS).

Figure 21: Two planetary nebulae, NGC 2438 (top) and NGC 2899 (bottom), as
they appear in the SHS and VPHAS+ surveys. The SHS images are shown in the
left-hand panels, with the VPHAS+ images to the right. The bands used to form
them are: NGC 2438 – SHS R/G/B = $H\alpha$/SR/SSS Bj, VPHAS+ R/G/B =
$H\alpha$/$r$/$i$, NGC 2899 – SHS R/G/B = $H\alpha$/SR/SSS Bj, VPHAS+ R/G/B =
$H\alpha$/$r$/$g$. The cut-out image dimensions are 300$\times$300 arcsec2 for
NGC 2438, and 200$\times$180 arcsec2 for NGC 2899.

##  7 Nebular astrophysics with VPHAS+ images

Just over a decade ago the SuperCOSMOS H$\alpha$ Survey (SHS, Parker et al
2005) had only just completed. This was the last survey using photographic
emulsions that the UK Schmidt Telescope undertook. The 3-hour narrowband
H$\alpha$ filter exposures reach a very similar limiting surface brightness to
the 2 minute exposures VPHAS+ is built around. Hence, the differences in
capability are not about sensitivity, as this is roughly the same in the two
surveys. Instead it is about the great improvement in dynamic range on
switching to digital detectors, the good seeing of the VST’s Paranal site, and
the added broad bands.

SHS, with its enormous 5-degree diameter field, has been comprehensively
trawled for southern planetary nebulae (the MASH catalogue, Parker et al 2006,
Miszalski et al 2008). The remaining discovery space for resolved nebulae is
expected to be at low surface brightnesses in locations of high stellar
density, and in the compact domain around and below the limits of the typical
spatial resolution of SHS ($\sim 0.5$ to 3.0 arcsec). Both these conditions
will most often be met in the Galactic Bulge, at a mean distance of $\sim$8
kpc. Data-taking in the Bulge and its maximally-dense star fields is planned
to begin in mid 2014.

Among planetary nebulae (PNe), small angular size is due either to great
distance or to youth – the study of either compact category provides exciting
possibilities. As well as the Bulge, the less-studied outer parts of the
Galactic Plane should be searched. In this respect, IPHAS, with its direct
view to the Galactic Anticentre is better positioned: the ongoing study of the
Anticentre PN population has revealed dozens of new candidates (Viironen et
al. 2009a), including the PN with the largest galactocentric distance to date
(20.8 $\pm$ 3.8 kpc, Viironen et al 2011). By following up such finds to
measure chemical abundances, crucial beacons are obtained for the study of the
Galactic abundance gradient and its much disputed flattening towards the
largest galactocentric radii. VPHAS+ completed the access to the outer Plane
over the longitude range $215^{\circ}<\ell<270^{\circ}$.

Data from both IPHAS and VPHAS+ can make fundamental contributions to the
study of very young PNe – particularly by helping to solve the two-decades-old
puzzle of how PNe already emerge with the observed wide variety of
morphologies (round, elliptical, bipolar, multipolar, point-symmetric, etc. –
see Sahai et al. 2011). What does this variety say about the properties of
their AGB progenitors? Detailed studies of objects in the phases preceding the
PN phase – AGB and post-AGB stars, proto-PNe, and transition or PN-nascent
objects – are underway (e.g. Sanchez Contreras & Sahai 2012). Superb imaging
capabilities like those of the VST, accessed via VPHAS+, will support this
work.

Indeed there is a serious paucity of very small PNe in the existing optical
catalogues: there are no PNe with angular extent less than 3 arcsec in the
MASH catalogue (out of 903 objects; Parker et al 2006), and only 8 PNe in the
catalogue by Tylenda et al (2003, 312 objects) in the size range 1.4 – 3
arcsec. There is just one with a confidently-measured diameter below 1 arcsec
in the larger Strasbourg Catalogue of PNe (1143 objects; Acker et al. 1994),
that happens to be a Bulge PN. IPHAS has demonstrated that extremely young
compact PNe can be reached (Viironen et al, 2009b), while Sabin et al (in
prep.) have found some 20 new PNe with diameters of 1-3 arcsec in by-eye
searchs of IPHAS image mosaics. Even smaller, but brighter, nebulae around
symbiotic stars of the dusty D subtype are emerging – the record so far being
IPHASJ193943.36+262933.1, a new D symbiotic star with an $H\alpha$ extent of
only 0.12 arcsec that has been confirmed via HST imaging and recently studied
with the 10.4m GTC telescope (Rodriguez Flores et al. 2014, submitted to A&A).

Apart from opening up new discoveries, a further benefit of good seeing is the
clearer view of nebular structure that it offers. This is nicely demonstrated
in fig 21. SHS and VPHAS+ detect the main features of the planetary nebulae
NGC 2438 and NGC 2899 to very similar depth – for example, the fainter outer
halo is just detected in both versions of NGC 2438. But, evidently, the VPHAS+
images better resolve the fine sculpting within both nebulae as a consequence
of the seeing FWHM being under a half that prevailing in SHS data. The
extended dynamic range of VPHAS+ helps in this respect, too, in that early
saturation also obliterates detail. This advantage is especially clear in the
images of NGC 2899, where the structure in the bright nebulous lobes is
preserved in VPHAS+, but is entirely bleached out in SHS. The more the level
of detail that can be picked out, the more certain and subtle morphological
classifications and interpretations can become.

The combination of good seeing and high dynamic range also makes the
separation of fainter stars from background nebulosity much easier. This
capability is critically important to the study of the young massive clusters,
still swathed in diffuse Hii emission, where the analysis of stellar content
is very much a focus of continuing research. For example, Feigelson et al
(2013) have offered a critique of the nuisance created by spatially-complex
nebulosity. The obvious answer to this and the problem of dust obscuration is
to turn to selection using NIR and X-ray data. Nevertheless the availability
of imaging data of the high quality seen in VPHAS+ data will make it possible
to extend SEDs for many more stars into the effective-temperature (and
reddening) sensitive optical domain. In addition, understanding the shaping of
the interstellar medium in star-forming environments remains an important part
of the picture (see e.g. Wright et al 2012 on proplyd-like structures in Cyg
OB2). The detail that the VST is capable of revealing both in obscuration and
ionised hydrogen in star-forming regions can be quite exquisite. Here, in fig
22, we illustrate this with an excerpt from VPHAS+ data on the Lagoon Nebula,
showing the fine tracing of the shapes of dark globules and eroding dusty
structures that is achieved.

Figure 22: A cut-out at full resolution from M8, the Lagoon Nebula. This is an
RGB image centred on RA 18 09 36 Dec -24 01 51 (J2000), and spanning
150$\times$150 arcsec2. The filters are combined such that R/G/B =
$H\alpha$/$i$/$r$. Figure 23: The central star of NGC 2899 reveals itself. The
top panel is a 1$\times$1 arcmin2 thumbnail of the centre of NGC 2899 as
imaged through the $u$ filter, while the bottom is the corresponding $r$
thumbnail. The white right-angled bars pick out the position of an extremely
blue, relatively faint star that is clearly present in all $u$ (and $g$)
exposures obtained, but is too faint for detection in $r$.

In planetary and other evolved-star nebulae it is of course important to
identify the ionising object. The search for missing PN central stars is a
quest that VPHAS+ can aid greatly through the provision of spatially well-
resolved $u$ and $g$ data. Indeed inspection of the data used to contruct fig
21 has revealed the probable central star of NGC 2899 for the first time. As
shown in fig 23, there is very evidently a third very-blue star just SW of the
pair of stars that have, in the past, been scrutinised as possible companions
to what is required to be an extremely hot ($T_{eff}>250,000$ K), but probably
faint central star (López et al 1991). This blue object was detected on the
night of 20th December 2012 at a provisional $u$ magnitude of 18.79 $\pm
0.02$. It fades through $g$ (19.36 $\pm 0.02$) to become undetected by the
pipeline, and scarcely visible to eye inspection, in $r$. Its coordinates are
RA 09 27 02.72 Dec -56 -06 22.9 (J2000), just 1.7 arcsec from the more
southerly of the pair of brighter stars examined before by López et al (1991).
Based on the $g$ magnitude and an inferred $V$ flux, we have determined the
central star’s effective temperature, via the well-established Zanstra method.
Using the reddening and integrated H$\alpha$ flux from López et al and Frew,
Bojičić & Parker (2013) respectively, we estimate $T_{\rm z,H}=215\pm 16$ kK.
This is cooler than the temperature given by López et al, based on the
’crossover’ method, but still extraordinarily hot for a central star well down
the white-dwarf cooling track.

It was one of the major science drivers for the merged VPHAS+ survey that $u$
data, supported by $g$, would result in the detection of a broad range of
intrinsically very blue objects – be they PN central stars, interacting
binaries or massive OB and Wolf-Rayet stars. An extreme example like NGC
2899’s central star provides the useful lesson that selection via the
$u-g,g-r$ colour-colour diagram would have failed to pick it out – because of
the non-detection in $r$. In a case like this, the $u,u-g$ colour-magnitude
diagram has to be examined, in tandem with the appropriate images.

##  8 VPHAS+ photometry as a reference set for variability studies

As the northern survey, IPHAS, has progressed over the decade since 2003,
there have been occasions on which it was possible to use the growing database
as a high-quality reference for checking transient reports – particularly of
novae. The most spectacular IPHAS example of this was the nova and variable,
V458 Vul (Wesson et al 2008; Rodriguez-Gil et al 2010) where the eruption
occurred a few months after obtaining H$\alpha$ images revealing a pre-
existing ionised nebula around the star. Indeed there have been several
instances in which photometry of the progenitor object has been extracted from
the IPHAS database and has been used to gain insight into the prior presence
or absence of line emission or to set constraints on likely extinction
(Steeghs et al 2007, Greimel et al 2012). Such opportunities will certainly
arise with VPHAS+ – and be richer given the five filters offered.

In the southern hemisphere novae will be more frequent, as will other
transient events. Furthermore responses to alerts, or the need to demonstrate
long-term flux variations, can bring into use repeats of observations made
necessary by initial quality-control failures. An example of this is provided
by Vink et al (2008) who used repeat IPHAS observations – taken on account of
poor observing conditions – to discuss the LBV candidacy of G79.29$+$0.46.
With the increased attention being given to the reporting and exploitation of
transient objects (including the forthcoming Gaia alerts programme), this use
of VPHAS+ will become more common.

##  9 Summary and concluding remarks

This paper has introduced and defined VPHAS$+$, the VST Photometric $H\alpha$
Survey of the Southern Galactic Plane and Bulge. The data taking, the
rationale behind it, the data processing and data quality have all been
described. The properties and limitations of the survey’s narrowband H$\alpha$
filter, NB-659, have been laid out and simulated in order to anticipate its
performance. In addition we have provided tables of the expected photometric
colours of normal solar-metallicity stars to aid the interpretation of the
survey’s characteristic photometric diagrams – most are to be found in the
Appendix where the effect of changing the adopted reddening law is
illustrated. The VPHAS+ H$\alpha$ filter transmission is redder, wider and
$\sim 20$ % higher-throughput than its IPHAS counterpart – a difference that
feeds through to noticeably different $(r-H\alpha)$ colours for M stars.

We have validated the photometry that is delivered by VST/OmegaCAM and
subsequently pipelined at CASU, using test data taken of a field for which
SDSS photometry is available. We find the agreement is satisfactory, with the
$g$, $r$ and $i$ band calibrations differing by between 0.01 to 0.05
magnitudes. However, for the time being, the pipeline calibration should be
regarded as provisional – it will undoubtedly improve. Examples of the
excellent imaging performance of the VST/OmegaCAM combination relative to
previous surveys have been provided, and we draw attention to the valuable
archival role this first digital survey can fulfill in supporting discoveries
of transient sources.

Exploitation of the survey is now beginning. The detection of a compact
ionized nebula around W26, the extreme M supergiant in Westerlund 1 has
already been published (Wright et al 2014). Applications have been made for
follow-up spectroscopy that will test the quality of selection of specialised
object types that VPHAS+ photometry makes possible. Progress is also being
made via direct analyses of the photometry. For example, Mohr-Smith et al (in
prep.) are conducting a search for OB stars in the vicinity of the massive
cluster Westerlund 2, and they are finding a close match between the
properties of known cluster O stars as derived from VPHAS+ data and those
inferred by Vargas Alvarez et al (2013, see also Drew et al 2013). This and
other early appraisals of the data indicate that VPHAS+ will be an excellent
vehicle for automated searches for reddened early-type stars. Kalari et al (in
prep.) are employing both narrow-band H$\alpha$ and the broadbands to measure
mass-accretion rates in pre-main-sequence stars: they are finding that
H$\alpha$ mass-accretion rates in T Tauri stars compare favourably to rates
determined from the $u$ band in the case of the Lagoon Nebula, NGC 6530.

As the calibration of the survey data improves, the measurement of accurate
integrated H$\alpha$ fluxes for many faint southern PNe and other extended
objects becomes possible, and will extend the work of Frew et al (2013, 2014).
In due course these fluxes can be compared with existing and also new radio
continuum fluxes coming on stream (see e.g. Norris et al 2011) in order to
determine reliable extinction values for many faint nebulae currently lacking
data. This technique has already been applied to the case of W26 in Westerlund
1 (Wright et al 2014).

When it becomes possible to cross-match VVV and VPHAS$+$ data, it will open up
the power of homogeneous photometric mapping of the central parts of the
Galactic Plane in up to 10 photometric bands spanning both the optical and the
near-infrared. Beyond the VVV sky area, there is a synergy to be exploited in
bringing VPHAS+ data together with those of the all-sky 2MASS survey
(Skrutskie et al 2006) and with the UKIDSS Galactic Plane Survey (Lucas et al
2008), in those parts of the first and third Galactic quadrants the latter has
covered. It is worth noting, however, that 2MASS alone is too shallow to link
effectively with VPHAS+ for sightlines where the integrated visual extinction
is less than $\sim 5$ magnitudes. This does mean that the longitude range
$230^{\circ}<\ell<300^{\circ}$, in particular, is presently lacking
sufficiently deep NIR photometry. In the longer term, many of the sources of
interest that VPHAS+ finds will benefit from accurate parallaxes and other
data from ESA’s Gaia mission – given the similar sensitivity limits reached.
Conversely in the meantime, VPHAS+ has already begun to assist ambitious wide-
field spectroscopy programmes such as the Gaia-ESO Survey (Gilmore et al 2012)
through the provision of the wide-field photometry needed for target selection
and field setup.

By the end of 2013, 25% of all observations making up the survey had been
obtained to the required quality, and in May 2013 a first release of single-
band catalogues was made to the ESO archive that contained roughly 10% of the
eventual total (based on data obtained prior to 15 October 2012). By design,
the characteristics of VPHAS+ are similar to those of the IPHAS and UVEX
Galactic plane survey pair in the north. In particular, the double-pass
strategy is shared, with the result that the majority of detected objects are
picked up and measured twice, with no more than $\sim$0.1 percent of objects
missed altogether. This feature has informed the way in which the IPHAS DR2
catalogue (Barentsen et al, 2014) has been constructed – and it is intended
that a first band-merged VPHAS+ catalogue, for public release, will be built
along analogous lines during the second half of 2014. This will incorporate
data from the first 3 seasons of VST observing, and give a complete
photometric account of the Galactic mid-plane. For ease of use, for every
detected source, the catalogue will provide a single recommended set of
magnitudes in up to 5 optical bands.

## Acknowledgments

This paper makes use of public survey data (programme 177.D-3023) obtained via
queue observing at the European Southern Observatory. In respect of the
H$\alpha$ filter, we would very much like to thank Bernard Muschielok for the
benefit of his expertise and support in connection with its laboratory
testing, and Jean-Louis Lizon for his steady hand in correcting some minor
surface defects. The referee of this paper is thanked for constructive
comments that improved its content.

This research made use of the AAVSO Photometric All-Sky Survey (APASS), funded
by the Robert Martin Ayers Sciences Fund. Many elements of the data analysis
contained in this work have been eased greatly by the TOPCAT package created
and maintained by Mark Taylor (Taylor, 2005). The pipeline reduction also
makes significant use of data from the Two Micron All Sky Survey (2MASS),
which is a joint project of the University of Massachusetts and the Infrared
Processing and Analysis Center/California Institute of Technology, funded by
NASA and the NSF.

JED and GB acknowledge the support of a grant from the Science & Technology
Facilities Council of the UK (STFC, ref ST/J001335/1). The research leading to
these results has also benefitted from funding from the European Research
Council under the European Union’s Seventh Framework Programme (FP/2007-2013)
/ ERC Grant Agreement n. 320964 (WDTracer). BTG was also supported in part by
the UK STFC (ST/I001719/1). RLMC and AMR acknowledge funding from the Spanish
AYA2007-66804 and AYA2012-35330 grants. HJF and MM-S both acknowledge STFC
postgraduate studentships. NJW is in receipt of a Royal Astronomical Society
Fellowship. RW acknowledges funding from the Marie Curie Actions of the
European Commission (FP7-COFUND).

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##  Appendix A Synthetic colour reddening tables

Synthetic colours for main sequence stars, computed as described in Section 4,
are tabulated in full in an online supplement for three representative
reddening laws ($R_{V}=2.5$, 3.1 and 3.8) and a range of reddenings
($A_{0}=0$, 2, 4, 6, 8, 10). The form of the reddening laws used is due to
Fitzpatrick & Massa (2007). As an example of the tables available we include
excerpts from the second and fifth tables that respectively provide
$R_{V}=3.1$ blue-filter and red-filter colours for B,A main-sequence stars.

Two further tables of synthetic colours are included in the supplement for K-M
giants that have been computed using P98 library spectra. Data are provided
for the $R=3.1$ mean Galactic law only, for the limited purposes of (a) giving
an impression of how these luminous red objects may contaminate $(u-g,g-r)$
diagrams at redder $(g-r)$ through $u$ red leak (b) enabling comparisons with
the M-giant spur commonly seen in $(r-H\alpha,r-i)$ colour-colour diagrams.

Table 6: VST/OmegaCAM synthetic colours for B,A main-sequence stars in the $(u-g),(g-r)$ plane reddened with an $R_{V}=3.1$ extinction law. (Full table online.) Spectral | $A_{0}=0$ | $A_{0}=2$ | $A_{0}=4$ | $A_{0}=6$ | $A_{0}=8$
---|---|---|---|---|---
Type | $(u-g)$ | $(g-r)$ | $(u-g)$ | $(g-r)$ | $(u-g)$ | $(g-r)$ | $(u-g)$ | $(g-r)$ | $(u-g)$ | $(g-r)$
$B0V$ | $-1.433$ | $-0.271$ | $-0.692$ | $0.529$ | $0.087$ | $1.301$ | $0.891$ | $2.050$ | $1.632$ | $2.777$
$B1V$ | $-1.324$ | $-0.240$ | $-0.584$ | $0.558$ | $0.195$ | $1.329$ | $0.995$ | $2.076$ | $1.719$ | $2.802$
$B2V$ | $-1.209$ | $-0.218$ | $-0.470$ | $0.579$ | $0.307$ | $1.350$ | $1.104$ | $2.096$ | $1.808$ | $2.821$
$B3V$ | $-1.053$ | $-0.186$ | $-0.315$ | $0.610$ | $0.460$ | $1.379$ | $1.250$ | $2.125$ | $1.923$ | $2.849$
$B5V$ | $-0.828$ | $-0.139$ | $-0.092$ | $0.655$ | $0.680$ | $1.423$ | $1.460$ | $2.166$ | $2.080$ | $2.890$
$B6V$ | $-0.728$ | $-0.121$ | $0.007$ | $0.672$ | $0.776$ | $1.439$ | $1.550$ | $2.182$ | $2.144$ | $2.905$
$B7V$ | $-0.580$ | $-0.100$ | $0.152$ | $0.692$ | $0.918$ | $1.458$ | $1.682$ | $2.200$ | $2.234$ | $2.922$
$B8V$ | $-0.388$ | $-0.076$ | $0.340$ | $0.714$ | $1.101$ | $1.478$ | $1.850$ | $2.219$ | $2.344$ | $2.940$
$B9V$ | $-0.198$ | $-0.046$ | $0.528$ | $0.742$ | $1.285$ | $1.504$ | $2.019$ | $2.244$ | $2.445$ | $2.964$
$A0V$ | $-0.053$ | $-0.005$ | $0.675$ | $0.780$ | $1.431$ | $1.540$ | $2.153$ | $2.277$ | $2.514$ | $2.995$
$A1V$ | $-0.019$ | $0.005$ | $0.709$ | $0.790$ | $1.464$ | $1.550$ | $2.181$ | $2.287$ | $2.525$ | $3.005$
$A2V$ | $0.021$ | $0.025$ | $0.749$ | $0.809$ | $1.505$ | $1.568$ | $2.217$ | $2.304$ | $2.538$ | $3.021$
$A3V$ | $0.038$ | $0.059$ | $0.771$ | $0.840$ | $1.531$ | $1.597$ | $2.241$ | $2.332$ | $2.541$ | $3.048$
$A5V$ | $0.067$ | $0.125$ | $0.805$ | $0.904$ | $1.567$ | $1.658$ | $2.269$ | $2.390$ | $2.523$ | $3.105$
$A7V$ | $0.044$ | $0.199$ | $0.788$ | $0.975$ | $1.554$ | $1.726$ | $2.252$ | $2.456$ | $2.474$ | $3.169$
Table 7: VST/OmegaCAM synthetic colours for B,A main-sequence stars in the $(r-i),(r-H\alpha)$ plane reddened with an $R_{V}=3.1$ extinction law. (Full table online.) Spectral | $A_{0}=0$ | $A_{0}=2$ | $A_{0}=4$ | $A_{0}=6$ | $A_{0}=8$ | $A_{0}=10$
---|---|---|---|---|---|---
Type | $(r-i)$ | $(r-H\alpha)$ | $(r-i)$ | $(r-H\alpha)$ | $(r-i)$ | $(r-H\alpha)$ | $(r-i)$ | $(r-H\alpha)$ | $(r-i)$ | $(r-H\alpha)$ | $(r-i)$ | $(r-H\alpha)$
$B0V$ | $-0.150$ | $0.054$ | $0.278$ | $0.198$ | $0.694$ | $0.316$ | $1.100$ | $0.409$ | $1.496$ | $0.478$ | $1.884$ | $0.526$
$B1V$ | $-0.136$ | $0.048$ | $0.291$ | $0.192$ | $0.708$ | $0.310$ | $1.114$ | $0.403$ | $1.510$ | $0.472$ | $1.898$ | $0.519$
$B2V$ | $-0.123$ | $0.045$ | $0.304$ | $0.188$ | $0.721$ | $0.306$ | $1.126$ | $0.398$ | $1.523$ | $0.466$ | $1.911$ | $0.513$
$B3V$ | $-0.104$ | $0.044$ | $0.323$ | $0.186$ | $0.740$ | $0.303$ | $1.145$ | $0.394$ | $1.541$ | $0.462$ | $1.929$ | $0.508$
$B5V$ | $-0.077$ | $0.039$ | $0.349$ | $0.180$ | $0.765$ | $0.295$ | $1.170$ | $0.386$ | $1.566$ | $0.452$ | $1.954$ | $0.497$
$B6V$ | $-0.068$ | $0.036$ | $0.358$ | $0.177$ | $0.774$ | $0.291$ | $1.179$ | $0.381$ | $1.575$ | $0.448$ | $1.963$ | $0.492$
$B7V$ | $-0.057$ | $0.029$ | $0.369$ | $0.170$ | $0.785$ | $0.284$ | $1.190$ | $0.374$ | $1.586$ | $0.440$ | $1.973$ | $0.484$
$B8V$ | $-0.045$ | $0.018$ | $0.382$ | $0.158$ | $0.797$ | $0.272$ | $1.202$ | $0.362$ | $1.598$ | $0.427$ | $1.985$ | $0.471$
$B9V$ | $-0.028$ | $0.006$ | $0.398$ | $0.145$ | $0.813$ | $0.259$ | $1.218$ | $0.348$ | $1.614$ | $0.413$ | $2.001$ | $0.456$
$A0V$ | $-0.009$ | $-0.005$ | $0.418$ | $0.133$ | $0.833$ | $0.246$ | $1.238$ | $0.334$ | $1.633$ | $0.399$ | $2.020$ | $0.441$
$A1V$ | $-0.003$ | $-0.003$ | $0.423$ | $0.135$ | $0.838$ | $0.248$ | $1.243$ | $0.335$ | $1.638$ | $0.399$ | $2.025$ | $0.442$
$A2V$ | $0.006$ | $-0.004$ | $0.432$ | $0.134$ | $0.847$ | $0.247$ | $1.251$ | $0.334$ | $1.646$ | $0.397$ | $2.033$ | $0.439$
$A3V$ | $0.021$ | $-0.008$ | $0.446$ | $0.130$ | $0.861$ | $0.241$ | $1.265$ | $0.328$ | $1.660$ | $0.391$ | $2.047$ | $0.432$
$A5V$ | $0.051$ | $0.005$ | $0.476$ | $0.141$ | $0.890$ | $0.250$ | $1.293$ | $0.335$ | $1.687$ | $0.396$ | $2.073$ | $0.436$
$A7V$ | $0.083$ | $0.027$ | $0.507$ | $0.160$ | $0.920$ | $0.268$ | $1.322$ | $0.350$ | $1.716$ | $0.410$ | $2.101$ | $0.448$