text:
\section{Introduction}

In the last twenty years increasing attention has been paid to the study of
dwarf galaxies in order to understand their crucial role in galaxy formation
and evolution. In hierarchical clustering theories (White \& Frenk 1991;
Kauffmann, White, \& Guiderdoni 1993) these systems can constitute the
building blocks from which larger systems have been created by merging,
while in monolithic collapse scenarios (Tinsley \& Gunn 1976; Tinsley 1980a)
they have been suggested to represent the debris of massive galaxies unable
to form stars until z$\sim$1. A population of newly star--forming dwarfs at
z$<$1 has been also invoked in some evolutionary models (Broadhurst, Ellis, \&
Shanks 1988;
Babul \& Ferguson 1996; see also Ellis 1997 for a review on the subject)
to reproduce the excess of faint blue galaxies observed in deep photometric
surveys, the most famous being the Hubble Deep Field (Williams et al. 1996).

Early--type dwarfs (dEs and dSphs) are gas poor and constituted by
intermediate and old stellar populations, while late--type dwarfs (dIrrs)
are gas rich and their light is dominated predominantly by very young stars
associated with bright HII regions, indicators of an ongoing star formation
process. The former kind of dwarfs show a smooth luminosity distribution
with a low surface brightness (like Sextants, NGC~147, Leo~I, Fornax, Sculptor,
etc.), while the latter group appears with a patchy intensity distribution
(NGC~6822, NGC~1569, IC~1613, etc.) and an intermediate surface brightness
which can become very high in the bright blue knots of blue compact dwarf
galaxies (hereinafter BCDGs).

At present it is not well understood if there is an evolutionary
interconnection between dIrr and dE galaxies
(see e.g. Gallagher 1998), as no consistent picture
of dwarf galaxy evolution has emerged yet. It has been suggested
(Davies \& Phillips 1989; and references therein) that the natural
evolution of dIrrs could be the condition of dEs through the phase of
BCDG. In this hypothesis, a strong starburst (intense star formation episode
concentrated in a very short time) can originate a galactic superwind
(Heckman 1995) with the mechanical energy supplied by stellar winds and
supernova explosions generated by newly formed massive stars. This wind
can blow out all the gas from a dIrr, due to its shallow potential well,
and transform the galaxy into a gas--poor dE. Indeed, narrowband H$\alpha$
images and X-ray maps show evidence of the existence of these superwinds in
some irregulars and BCDGs, as for example NGC~1569 (Heckman 1995), NGC~1705
(Meurer et al. 1992), IZw18 (Martin 1996; Petrosian et al. 1997), etc.
On the other hand, taking into account observed chemical, photometric and
kinematic properties of both dwarf irregulars and ellipticals, it seems quite
hard to find an efficient mechanism to transform a late-type dIrr into a dE
(Jerjen \& Binggeli 1997). The derivation of the star formation history of
dwarf irregular galaxies and the corresponding identification of objects with
a starbursting regime become thus of primary importance to gain an insight
into the nature and evolution of dwarf galaxies in general.

IZw18 (also Mrk 116 or UGCA 166) is possibly the BCDG with the most striking
properties. At a recession velocity of 745$\,\pm\,$3 km/sec (Dufour, Esteban,
\& Casta\~{n}eda 1996a), corresponding to a distance of 10 Mpc (H$_0\,$= 75
km\,sec$^{-1}$Mpc$^{-1}$), this system shows very blue colors. The most
recent estimates on emission-line corrected broadband images
give U--B=\,--\,0.88 and B--V=\,--\,0.03 (Van Zee et
al. 1998). These colors are indicative of a very young
population, but do not exclude an underlying older one.
The total mass of IZw18 from the rotation curve at a radius of
10$^{\prime\prime}\!$--12$^{\prime\prime}$ is estimated  to be
$\sim 10^8$ M$_{\odot}$ (e.g. Davidson \& Kinman 1985;
Petrosian et al. 1997; Van Zee et al. 1998).
A large amount of neutral gas is detected all around the system,
totalling $\sim 7\times10^7$ M$_{\odot}$. This
corresponds to $\sim$70$\%$ of the total mass, but only $ 10^7$
M$_{\odot}$ of HI are associated with the optical part of the galaxy
(e.g. Lequeux \&
Viallefond 1980; Van Zee et al. 1998). When discovered by Zwicky (1966),
IZw18 was described as ``two galaxies separated by 5\farcs6 and
interconnected by a narrow luminous bridge'', surrounded by two ``very faint
flares'' at 24$^{\prime\prime}$ northwest. More recent CCD ground--based
images (Davidson, Kinman, \& Friedman 1989; Dufour \& Hester 1990; hereinafter
respectively DKF89 and DH90) have revealed a more complex structure: the
{\it two galaxies} are in fact two
star--forming regions of the same galaxy (usually indicated as NW and SE
components), while the {\it two flares} are just the most prominent of a few
nebulosities surrounding IZw18. These minor systems are roughly aligned
toward the northwest and were initially believed at the same distance,
but now we know from spectroscopic studies that only one component
(referred to as component C in DKF89) is at the same distance
as IZw18
and is physically associated with the main body (Dufour et al. 1996a;
Petrosian et al. 1997; Van Zee et al. 1998). The
other diffuse objects have been recognized as background galaxies (see Fig.~1a
of Dufour et al. 1996b, D96, for an overview of the whole IZw18 system). In
the following we will refer to IZw18 and to component C respectively as main
body and secondary body (or companion) of IZw18.
Both systems have been resolved into single stars
for the first time only with HST/WFPC2 by Hunter \& Thronson (1995,
hereinafter HT95) and D96.
Indeed, IZw18 and its companion
irregular galaxy are currently one of the most distant systems ever
resolved into stars.

This apparently insignificant BCDG became famous right after its discovery,
when Searle \& Sargent (1972) measured from its emission-line spectrum an
oxygen abundance [O/H]= --1.14, corresponding to only 7$\%$ of the solar value
and indicating a quite unprocessed gas content. Furthermore, the first studies
on its color and composition (Sargent \& Searle 1970; Searle \& Sargent 1972;
Searle, Sargent, \& Bagnuolo 1973) already emphasized  a current star
formation rate (SFR) much
higher than the mean value in the past. All these observational evidences
brought to the formulation of the basic question on the nature of IZw18: is it
a young galaxy which is presently experiencing its first burst of star
formation, or is it an old system which has already formed stars in the past
in at least another episode of star formation ?

Subsequent spectroscopic studies in IZw18 (Lequeux et al. 1979; French 1980;
Kinman \& Davidson 1981; Davidson \& Kinman 1985; Dufour, Garnett, \&
Shields 1988; Garnett 1989, 1990; Pagel et al. 1992; Skillman \& Kennicutt
1993; Kunth et al. 1994; Stasi\'nska \& Leitherer 1996; Garnett et al. 1997;
Izotov \& Thuan 1998) have confirmed its extreme metal deficiency, around
1/30--1/50 of Z$_{\odot}$. Despite many efforts to detect other galaxies
with very low metallicity (Terlevich, Skillman, \& Terlevich 1995), IZw18 still
remains the galaxy with the lowest metal and helium content known so far. This
makes the system a fundamental point in the derivation of the primordial
helium abundance (Izotov, Thuan, \& Lipovetsky 1994, 1997; Olive, Steigman,
\& Skillman 1997; Izotov \& Thuan 1998) and in the study of the properties
of chemically unevolved galaxies. However, there are several observational
indications that IZw18 is not a primordial galaxy, for instance the presence
of relatively high C/O and N/O abundance ratios justified only with an earlier
population of low and intermediate mass stars (Dufour et al. 1988; Garnett et
al. 1997) and the photometric evidence of an underlying red stellar population,
both from surface photometry of the whole galaxy in the NIR (Thuan 1983) and
from photometry of single stars in the optical bands (HT95, D96).

For nearby galaxies the safest determination of their SF
history is obtained resolving their stellar population into single stars
and inferring their SFR and initial mass function (IMF) with the synthetic
color--magnitude diagram (CMD) method. This method was first developed for
dwarf irregular galaxies in the Local Group observed with ground-based
telescopes (Ferraro et al. 1989; Tosi et al. 1991, hereinafter TGMF; Greggio et
al. 1993, GMTF; Marconi et al. 1995, MTGF) and has now been updated (Greggio et
al. 1998, hereinafter G98) for an optimized application to galaxies observed
with HST. A procedure for the comparison between observed and synthetic CMDs
has been developed also by Tolstoy \& Saha (1996), who have introduced the
concept of Bayesian inference to give the relative likelihood of different
models to constitute a suitable representation of the data. Gallart et al.
(1996) also follow a similar approach and have introduced in the method the
concept of metallicity evolution following a given law for the chemical
enrichment of the interstellar medium. Here we have applied the synthetic CMD
method to IZw18 using the HST archive data from HT95 and D96.

The data reduction is described in Sect.2 and the resulting CMDs and LFs in
Sect.3. The method and a description of the comparison of these
data with theoretical synthetic CMDs and LFs are given in Sect.4, with
the resulting conclusions on the recent evolution of IZw18 and its companion.
An overall discussion of these conclusions in the
framework of the current common knowledge on this galaxy is finally given
in Sect.5.

\section{Observations and Data Reduction}

IZw18 has been observed with different instruments on board of
HST\footnote{Observations with the NASA/ESA Hubble
Space Telescope, are obtained at the Space Telescope Science Institute,
which is operated by AURA for NASA under contract NAS5-26555}
by different investigators. For our purposes we have retrieved from the
HST archive and re-reduced all the HST/WFPC2 images available at January
1st, 1998.

\subsection{The data}

Two sets of deep exposures were taken in
November 1994 and a third set of shorter ones in March 1995.
In the first set of data (PI: Hunter, GO-5309, November 1994) IZw18 was
centered on the PC CCD, with an effective plate scale of
0\farcs045 pixel$^{-1}$ and a field of view corresponding to
36$^{\prime\prime} \times$ 36$^{\prime\prime}$. The target
was observed in the three broadband filters F336W, F555W and F814W
(similar to the standard ground-based broadbands U, V and I),
and in the two narrowband filters F469N and F656N (sampling the nebular lines
HeII $\lambda$4686 and H$\alpha$ $\lambda$6563), in order to map
the ionized gas and WR stars. Results from
this set of data are presented in HT95.

The second set of exposures (PI: Dufour, GO-5434, November 1994) consists
of deep frames of IZw18 and its companion system on the WF3 CCD, with a
plate scale and a field of view of 0\farcs1 and 80$^{\prime\prime}
\times$ 80$^{\prime\prime}$ respectively. The frames are available in
the three broadband
filters F450W, F555W and F702W (corresponding indicatively to B, V and
R), and in the two narrowband filters F502N and F658N, mapping the
two nebular lines [OIII] $\lambda$5007
and [NII] $\lambda$6583. Photometric results from this set of data are
presented by D96.

Finally, the third set of images (PI: Dufour, GO-5434, March 1995), are
in the three broadband filters F439W, F555W and F675W (the F439W is a filter
in the B band region narrower than the F450W; the F675W is a filter in
the R band region narrower than the F702W).

A complete summary of all the data available for IZw18 as observed with the
HST/WFPC2 is presented in Table 1, where we have indicated
the filter (column 1), the WFPC2 camera where the target was centered on
(column 2), the principal investigator (column 3),  the epoch of observation
(column 4), the integration time in seconds for each single exposure (column
5) and the image root names (column 6).

We actually used only a subset of all the data available on IZw18,
as indicated in Table 1 by an asterisk near the image name. In fact, we were
interested in obtaining the color-magnitude diagrams V, B--V and V, V--I for
IZw18 with the highest possible resolution (i.e. from the PC camera)
and the color-magnitude diagram V, B--V of its companion on WF3.
The observations in the narrowband filters were used to take into account
the contribution of the ionized gas in the different bands.

\subsection{Photometric reduction}

We reduced the data applying all the corrections required
to minimize photometric uncertainties and to achieve the highest photometric
accuracy. All the reductions were performed in the IRAF
\footnote {IRAF is distributed by the National Optical Astronomy Observatories,
which is operated by the Association of Universities for Research in Astronomy,
Inc., under cooperative agreement with the National Science Foundation.}
environment.

For each single exposure we corrected warm pixels and
flagged hot pixels in the data quality files using the IRAF/STSDAS task
{\it warmpix}. To restore the correct relative count numbers
between pixels at different positions on the CCD,
we performed geometric distorsion and CTE corrections on each single
frame by applying respectively the correcting image f1k1552bu.r9h available
from the Archive (see Leitherer 1995 for more details) and the linear-ramp
image with the appropriate value depending on the background as indicated in
Holtzman et al. (1995a, b; hereinafter H95a and H95b).

Multiple frames through each filter in each dataset were simultaneously
co-added and cosmic-ray removed. We also removed where possible the
contribution of ionized gas from broadband images, proceeding as follows:
after having adequately smoothed narrowband images to eliminate pointlike
sources from gas maps, we calculated with SYNPHOT the percentage of
flux detected in each narrowband filter and subtracted it from each broadband
image.

The pre-reduction provided the following data useful for our purposes:
for IZw18 and its companion on the WF3 camera, four deep
frames in the F555W, F450W, F502N and F658N filters (total integration time
of 4,600 seconds for each filter); for IZw18 on the PC camera we have two
images in the F555W filter, the deepest one obtained from Hunter's frames
(for a total integration time of 6,600 seconds) and the shorter one from
Dufour's data (1,200 seconds in total). We preferred not to
combine the two final PC F555W frames,
because of a great rotational
displacement of one frame with respect to the other, which would have implied
a repixelization and data manipulation.
Also available are: one deep image
in the F814W band (6,600 seconds), a less deep frame in the F439W band (2,000
seconds) and a quite deep frame in F656N filter (total integration time of
4,200 seconds).  In Figs~\ref{VimagePC} and \ref{VimageWF3}
we show the deepest WFPC2 images in the F555W filter respectively for IZw18
on the PC camera and for its companion galaxy on the WF3 detector:
both systems are well resolved into single stars. It is also possible to
distinguish HII regions in the SE part of IZw18 as well as the NW cluster,
while 2 bright star clusters are evident in the center and in the NW part
of the companion system.

The photometric reduction of the frames was performed using the
DAOPHOT package in IRAF for PSF-fitting photometry in crowded fields on both
original and gas-subtracted broadband images. First we applied the automatic
star detection routine {\it daofind} to the deepest F555W frame
(detection threshold at 4$\sigma$ above the local background level), and
then we performed an accurate inspection by eye of each single detected object
to reject any feature misinterpreted as star by the routine
(namely, nuclei of faint galaxies, PSF tendrils, noise spikes, etc.).
The identification of stars in the other filters was then forced assuming the
final positions of the stars detected in the F555W deeper image as input
coordinates for the starting centering and aperture photometry in the new
frame. We also tried to force
the photometry from the deepest F555W image to the rotated F439W image,
taking into account the coordinate transformation, but we obtained
systematically higher photometric errors, introduced by transformation
uncertainties. We therefore decided to couple the F439W image with the
shallower but unrotated F555W one.

In spite of the performed centering, forcing the photometry in the second
band leads to some mismatches.
These have been identified by plotting the distance
between the centers in the two images as function of the
F555W instrumental magnitude. To overcome this problem, we used the two
routines {\it daomatch} and {\it daomaster}, kindly made available by P.
B. Stetson, to finally match the coordinate lists in the two coupled filters.

In order to obtain the best PSF for our frames we
experimented three different methods: a) we used some well isolated stars
(3 or 4) in each frame to build the observed PSF, b) we ran the {\it Tiny Tim}
software (Krist \& Hook 1996) to obtain a theoretical PSF and, c) we made
use of the new tool of the WFPC2 PSF library to get empirical PSFs. The
different PSFs considered give quite similar results for the PC camera;
eventually we preferred to adopt the observed PSF since it takes into
account all technical conditions occurring at the epoch of data
acquisition (real focus value, thermal breathing, etc.). For the WF3 camera
theoretical PSFs seem instead to work better, since in this case we have to
deal with a more dramatic undersampling of the observed/empirical PSFs which
introduces a higher photometric uncertainty.

Once all the stars were measured with the PSF--fitting, those with a disturbed
image (as indicated by the two image--peculiarity indices $\chi^2$ and {\it
sharpness}) were identified and rejected. The index $\chi^2$ gives essentially
the ratio of the observed pixel--to--pixel scatter in the fitting residuals to
the expected scatter based on the values of the detector characteristics
(readout noise and gain). The {\it sharpness} is related to the intrinsic
angular size of the astronomical object. We removed all the objects with
$\chi^2\!>\,$3 and {\it sharpness} lower than --1 or larger than +1, i.e.
objects with size smaller than a star (like cosmic rays or image defects)
or objects too extended (like blends or semi-resolved star clusters, HII
regions and galaxies). We also checked individually all the stars in critical
positions in our reference CMDs (i.e. the CMD of stars with photometric error
smaller than 0.2 mag in both filters), like
very bright blue and red stars (possible blends, unresolved stellar clusters,
or unidentified cosmic rays), very faint objects (possible
peaks of noise in the forced filters, or residuals of cosmic rays detections).
Finally, we checked accurately all the red stars, which are particularly
important to discriminate among different SF histories.

\subsection{Calibration}

The instrumental magnitudes obtained with the PSF--fitting technique were then
converted into calibrated magnitudes following the prescription of H95a and
H95b. Since in H95a and H95b the standard calibration is given for an aperture
of 0\farcs5, we transformed the instrumental magnitudes on an
aperture radius of 2 pixels into the corresponding ones on an aperture radius
of 0\farcs5 (11 pixels for PC and 5 pixels for WF3) by calculating
the aperture correction.

This turned out to be a very delicate step, due to the small number of isolated
stars, suitable for aperture photometry, in the field of
IZw18. The derived aperture corrections strongly depend on
the choice of the adopted stars, and the small number of good
templates leads to a large statistical uncertainty. Figure~\ref{apcorr}
illustrates this point for the deepest V and I frames.
For all the isolated stars in each frame (5--15 objects) we measured the
instrumental magnitudes based on aperture photometry
with different radii. For each star, the aperture corrections (i.e.
the difference between the magnitude at a certain aperture radius and
the PSF-fitting magnitude) are shown as open circles for different
values of the aperture radius. The full dots represent the
correction averaged over all the measured stars, with a
2$\sigma$ rejection algorithm, and the vertical bars show the
corresponding 2$\sigma$ ranges. To these mean
aperture corrections we added the encircled energy corrections as
indicated by H95b; the obtained values corresponding to the conventional
radius of 0\farcs5 are indicated by crosses.

As apparent from Fig.~\ref{apcorr}, the uncertainty in the calibration
increases with the aperture radius, with more points deviating
from the mean value, as it becomes increasingly difficult to find no
defects in the outer pixels which are progressively included in the
calculation. However, once averaged and corrected for the encircled
energy, the aperture correction turns out fairly constant. The mean of
the aperture correction values over the
considered range of aperture radii is indicated in Fig.~\ref{apcorr}
by the horizontal thick line, which corresponds to
corrections of $-0.38$ mag and $-0.54$ mag in the $V$ and $I$ bands
respectively. These values are almost identical to those given by
H95b, which is encouraging. We thus used directly the
H95b corrections for our deepest F555W and F814W frames.
For homogeneity, we extended the use of H95b aperture corrections to
the other frames of IZw18 on the PC camera,  as
well as to the frames of the companion galaxy
on WF3, where the aperture corrections are much more difficult to determine
empirically.

We finally corrected our measures for the gain factor, the WF3 normalization
and the contamination effect, where necessary, and applied the zero
points to scale the photometry into the WFPC2 synthetic system relative to
Vega (Table 9 of H95a). When the value of one of these correcting parameters
was not available (as for the F450W filter)  we took it from public images
of calibration programs (e.g. the aperture correction for F450W on WF3) or
assumed it from the corresponding value of similar filters (F439W instead of
F450W).

\subsection{Completeness analysis}

One of the larger uncertainties affecting galaxy photometry is due to
crowding. Thus we carried out an accurate completeness analysis using
the DAOPHOT routine {\it addstar}.
For each frame we performed a series of tests, by adding each time $\sim$10\%
of the stars detected in each half-magnitude bin on the original
image. We then performed a new photometric reduction of the frame using
the same procedure applied to the original frame, and considering the
same rejection criteria for spurious objects. We then checked how many added
stars were lost either because not detected by the automatic routine or because
recovered with a large mag difference $\Delta$m which makes them migrate
to another magnitude bin. We eventually estimated the completeness
factors by averaging the results obtained repeating the test
10--20 times on each frame. Due to the uneven distribution of stars on the
images, fainter objects are more easily recovered in the outer, less crowded
regions. In order to derive the average completeness factors affecting our CMDs
we constrained the {\it Addstar} routine to put artificial stars in the frame
regions where the galaxy is actually located.

For every pair of filters used to construct the CMD, we performed
the completeness analysis independently on each frame. Since we forced the
stellar detection in the shallower frame, the overall completeness factor is
actually that corresponding to the shallower filter: in practice, either
the B or the I frame. Table 2 shows the completeness factors (percentage of
recovered artificial stars) and the
corresponding uncertainties as a function of magnitude for all the frames
considered in our photometric analysis. The listed values take also into
account the star selection with photometric error smaller than 0.2 mag
described in the next section. The completeness factors are averages over
the whole galaxy. Clearly, incompleteness
can be quite different from one region to the other, depending on crowding.
For instance in the most crowded zone in the NW cluster, incompleteness
is so dramatic that we recover only $\sim$20\% of the artificial stars already
at the 23 magnitude. HII regions and the rich star clusters are unresolved
in our images and are therefore to be considered fully incomplete. Although
it could be more appropriate to apply different completeness factors in
different regions, when computing our theoretical simulations we apply only
the average completeness factors in Table 2. Indeed the resolved objects are
not numerous enough to allow us to simulate individual different regions.

The completeness tests allowed us also to evaluate the influence of
blending in our photometry and the goodness of the applied rejection criteria
($\chi^2$, sharpness, error, etc.).
In the deepest V and I images of IZw18 on the PC detector the rejection
criteria seem to affect more intermediate/faint magnitudes
(24$\la$V$\la$27 and 23$\la$I$\la$25), while in the shallower B and V frames
they remove more objects at brighter magnitudes (22$\la$V$\la$24.5 and
22$\la$B$\la$24). A large fraction of the artificial stars rejected for
$\chi^2\!>\,$3 and {\it sharpness} outside the range from $-$1 to 1
(actually a few objects, usually
rejected for the $\chi^2$ parameter) were recognizable as blends: we have been
able to actually see the companion star for $\sim$50\% of the objects rejected
in the B image, for $\sim$65\% in both the V images and $\sim$75\% in the I
image. For IZw18's secondary body on WF3 there were very few objects discarded,
and they were all blends in both the V and B frames. This result confirms the
need of applying these criteria in order to remove from the CMDs spurious
objects due to blends.

We also looked at the effect of blending on the derived stellar
magnitude, by selecting the artificial stars which were
recovered with a magnitude brighter than $\Delta$m=0.25 mag with respect to
the input value. This happened for 4.7\% and 4.0\% of the stars added
in the deeper V and I frames of IZw18; for 2.9\% and 5.2\% in the shallower
V and B images, and around 3.5\% in both the V and B frames for the companion
system on WF3.
These values give an estimate of the frequency of cases in
which blending affects the photometry more than our allowed photometric
error in the CMDs ($\sigma_{DAO}\!<\,$0.2 mag in both filters).
We conclude that
our data are affected by blending at an average level of $\sim$4\%,
which we consider negligible for the interpretation of the CMD with
the simulation procedure.

\subsection{Photometric errors}

Figures~\ref{errPC} and \ref{errWF3} show the behavior of the photometric error
$\sigma_{DAO}$ (estimated by DAOPHOT) as a function of magnitude, respectively
for IZw18 in each combined PC image and for the companion galaxy
in each combined WF3 frame. In these plots we considered all the stars fitted
in both coupled filters: respectively 568 and 321 objects for the V vs V--I
and V vs B--V diagrams relative to IZw18, and 117 stars for the V vs B--V
diagram of the companion.

When the rejection criteria of $\chi^2$ and {\it sharpness} are applied,
most of the points with high $\sigma_{DAO}$ at bright magnitudes disappear,
since they are most probably small unresolved stellar associations,
HII regions or blends.

For the theoretical interpretation we will restrict our CMD to objects
with photometric error smaller than $\sim$ 0.2 mag. We notice that
$\sigma_{DAO}$ remains below
0.2 mag down to m$\sim$26 in each deeper frame of both IZw18 and the companion,
and down to m$\sim$25 in the less deep ones of IZw18.
However, $\sigma_{DAO}$ may underestimate by $\sim$20\% the
total photometric error (Stetson \& Harris 1988) and we have therefore
derived an independent estimate of the latter using the outcome of
the completeness tests, and looking at the amplitude of the
difference $\Delta$m between the assigned and recovered magnitudes
of the artificial stars.  It turned out that for the companion of IZw18,
the upper envelope of the distribution of the $\Delta$m
coincides with that of the observed $\sigma_{DAO}$, indicating that in this
case the errors are well estimated. For the PC frames with IZw18, instead,
$\sigma_{DAO}$ underestimates by more than 20\% the actual error, especially
at the brightest magnitudes. The larger error estimates $\Delta$m have
therefore been adopted in the CMD simulations described in Sect.4.

\subsection{Comparison with previous photometric analyses}

We have compared the photometric results obtained for IZw18 in the deep
F555W and F814W images on the PC camera with the results of HT95. The
comparison of our CMD (see Fig.~\ref{VIcmdPC}) with that published by HT95
(see their Fig.~5) shows that we reach a fainter limiting magnitude. Our
diagram contains many more blue and red stars fainter than V$\sim$26.
We have carefully checked the faintest objects to reject any uncertain
detection with ambiguous shape or profile, but most of
them remain and are localized in the outer and less crowded part of the
SE star-forming region where the photometry of fainter objects is easier.
A possible reason for our fainter limiting magnitude could be our
forcing the photometry in the shallower filter.

We re-measured all the stars listed in Table 2 of HT95, adopting their x and y
coordinates, in order to directly compare our mags with theirs: despite
the same photometric package (DAOPHOT) and the same parameters
for the photometric conversion of the instrumental magnitudes into the
calibrated ones (zero points, gain, contamination, etc.),
we found a shift in the zero points of $\sim$0.22 in F555W and $\sim$0.14
in F814W, in the sense that we are
brighter and bluer than HT95. This can be seen in Fig.~\ref{photcomp}.
The difference may arise from the use of gas subtracted images or from
the PSFs and aperture corrections. However, we have verified that there is
no significant
difference in the photometric reduction of images with and without gas
subtraction (except for some peculiar objects), and that there is no shift
in the zero points resulting from the use of PSFs obtained with different
techniques (as described in \S~2.2). We thus believe that the offset is
due to the difficulty in estimating the aperture corrections
from the stars in the images of this very problematic field.
As described in \S~2.3, there are too few isolated and reliable stars and the
dependence of the photometric correction on the chosen stars, shown in
Fig.~\ref{apcorr},
may well lead to offsets of some tenths of mag in the calibration.

We have compared also our CMD in the F450W and F555W filters for IZw18's
companion in the WF3 camera (see Fig.~\ref{BVcmdWF3}) with
the corresponding diagram published by D96 in their Fig.~3, accounting
for the reddening correction. The
general aspect of the two distributions is quite similar, despite
some differences. In our CMD there are less faint objects with unphysical
color (B--V)$\la-$1, probably thanks to the task for
masking hot and warm pixels (that could be mistaken as faint stars), which
has become available after the photometric reduction by D96.
There is also an offset in the zero points of both filters, with our
distribution being $\sim$0.5 mag brighter than D96 in V and in B.
D96 worked on gas-subtracted and rebinned images in order to construct
an artificial PSF better sampled than the observed one, while we preferred
to work on the original frames (but, again, we found no differences at all
using the gas subtracted images) and to use a PSF simulated
with the {\it Tiny Tim} software. As for the HT95 data, we
attribute the photometric offsets to uncertainties in the estimate of
the aperture corrections, enhanced in this case by the
undersampling of the WF3 camera.

For what concerns the CMD in the F439W and F555W filters of IZw18 on the PC
camera, the archival data  have not been published yet. We can therefore
compare only indirectly our resulting CMD (see our Fig.~\ref{BVcmdPC}) with
that derived from the F450W and F555W filters on the WF3 detector and
plotted in D96's Fig.~2. Both diagrams show a fairly large sequence
and are very similar to each other. When accounting for the reddening,
our edge of the blue plume
(at V$\,\simeq\,$22.5), the faint limiting magnitude (at V$\,\simeq\,$27),
and the average color of the star distribution around (B--V)$\,\simeq\,$0
seem consistent with those in D96.

\section{Observed color-magnitude diagrams and luminosity functions}

\subsection{IZw18 Main Body}

Figures~\ref{VIcmdPC} and \ref{BVcmdPC} show the CMDs derived from the
PC images of IZw18 in the HST
F555W, F814W and F439W bands, that in the following will be referred
to as the V vs V--I  and the V vs B--V CMDs respectively.
In panel a) of each figure we plotted
all the objects measured in both filters with a $\chi^2\!<\,$3 and
--1$<\!sharpness\!<$1 (respectively 444 in the V vs  V--I and 267 in the V
vs B--V), while in panel b) we show only the stars with a photometric
error $\sigma_{DAO}\!<\,$0.2 in both bands, after removing
spurious detections (247 objects remain in the red CMD and 106 in
the blue one). The main features
of the complete diagram remain unaltered after the $\sigma_{DAO}\!<\,$0.2
selection. Thus we used the CMDs in panels b) as reference diagrams
for the theoretical simulations, as they have a more reliable photometry.

Since this target is located at high galactic latitude (b=+45$^{\circ}$), its
CMDs do not suffer from significant contamination from foreground stars
belonging to our Galaxy, and it is not necessary to consider this factor of
uncertainty in our simulations or to correct for it our observed CMDs.

A first analysis of the observed V vs V--I diagram of IZw18 (panel b) of
Fig.~\ref{VIcmdPC}) shows that we
reached a limiting magnitude of V$\,\simeq\,$26.5 which goes down to
V$\,\simeq\,$27 for objects with the reddest V--I color. From a morphological
viewpoint, in this CMD we can easily distinguish the typical blue plume
observed in ground-based observations of Local Group irregulars: this plume
is populated both by main-sequence (MS) stars and by stars at the hot edge of
the post-MS evolutionary phases. For IZw18 the plume extends up to
V$\,\simeq\,$22 and has a median color V--I$\,\simeq\,$0, indicative of a very
low reddening. This is in agreement with the most recent values proposed for
this parameter, e.g. E(B--V)=0.04 in HT95.
We can also notice the presence in the CMD of several bright supergiants
with a wide spread of colors from blue to red,
and some faint red stars. As described in the previous sections, they
all turned out to be real stars after a detailed analysis of their
shape and profile.
All the bright blue and red stars are concentrated in the innermost and more
crowded regions, particularly in the NW component of IZw18.

All the HII regions recognized by HT95 (as well as many stars in the badly
resolved NW cluster), were automatically removed from our V, V--I diagram
because they turned out to have $\chi^2\!>\,$3 and/or $sharpness$ outside the
range [--1,1]\,. Also the bright star clusters and associations are rejected
for the same reason.
This implies that our CMD is not sampling extremely young stars. The
percentage of flux discarded with this procedure corresponds roughly to
3.5\% (2\%) of the total flux of the whole galaxy in the V (I) filter, while
it is $\sim$40\% of the total flux sampled by the resolved stars in both bands.

An inspection of the V vs B--V diagram shows that also in this case we reached
the limiting magnitude V$\,\simeq\,$26. We have however a shallower cut-off as
the B--V color becomes redder: this is due to the shorter integration time and
to the lower sensitivity of the F439W filter with respect to the F814W. Again
in this CMD we can recognize the typical blue plume of the MS and post-MS
stars, with a median color of B--V$\,\simeq\,$0, and an upper
brightness limit of V$\,\simeq\,$22.5. Also in this case we are not retaining
the HII regions of HT95 and the star clusters,
the total flux of the galaxy lost with the rejection criteria being
1\% (0.5\%) in the V (B) band, again $\sim$40\% of the light in the
resolved stars.

To estimate the masses of the stars visible in the CMDs and the
corresponding lookback times, we have converted stellar evolutionary tracks
into the observational plane, and superimposed them on the observed CMDs.
In Fig.~\ref{tracksPC} we show the Padova tracks with Z=0.0004
(Fagotto et al. 1994) converted to the V vs V--I (panel a) and V vs
B--V (panel b) plane, having adopted a distance modulus of (m--M)$_0$=
30 and a reddening value E(B--V)=0.04. The V vs V--I diagram
shows that in IZw18 we have detected MS stars
with masses higher than 12 M$_{\odot}$ (corresponding to lifetimes younger
than $\sim$ 20 Myr) and blue-loop stars
with masses down to $\sim$3--4 M$_{\odot}$ (thus with ages up to
$\sim$0.2 Gyr). The faintest clump of red objects in Fig.~\ref{tracksPC}
can be populated by (red) core helium burning, asymptotic giant branch (AGB)
stars and bright red giant branch (RGB) stars,
whose masses can be in principle as low as $\sim\,$1 M$_{\odot}$,
extending the lookback time up to several Gyr. However, given the large
photometric error, it is not possible to estimate precisely the mass, and
therefore the age, of these faint stars. In the following we will
conservatively assume a lookback time of 1 Gyr.
In the V vs B--V diagram we see objects with mass larger than 12 M$_{\odot}$
in the MS stage ($\tau\la\,$20 Myr) and larger than 7 M$_{\odot}$ in
the post-MS phase, thus with ages less than $\sim$ 50 Myr. We will use the
V vs V--I diagram to infer the SF history of IZw18 over the last
$\sim$ 1 Gyr, while the V vs B--V diagram will be used as a further
check over the last $\sim$ 50 Myr.

In Fig.~\ref{PCLFs} we plot the differential luminosity functions (LFs) of
all the stars with $\sigma_{DAO}\!<\,$0.2 in both filters present in
the V vs V--I and the V vs B--V diagrams (panels a) and b) respectively).
We can see in both cases a rather smooth trend. The LFs will be used in the
simulations to check the consistency between models and observations.

It is clear from Fig.~\ref{tracksPC} that the blue plume of IZw18 is
populated by stars both on the MS and at the hot edge of the blue loop
evolutionary phase and that no safe criterion can be found to separate
the two different populations. For this reason we do not even attempt to
derive a MS-LF, which would be inevitably affected by too large uncertainties
to be of any use.

Also the derivation of the slope of the LF may turn out too uncertain, once
we consider that, as listed in Table 2, the data start to be
incomplete already at the brightest mags and significantly incomplete at
V=24. For mere sake of comparison with other galaxies, and warning that these
values should only be taken as indicative, we have nonetheless
computed the slope by means of a maximum likelihood fitting on the deeper
V data. Down to V=23,
where the data are almost complete, but only very few stars are present,
 $\Delta$log\,N/$\Delta$V = 1.28 $\pm$ 0.04; at V=24, where completeness
is 85\%, the slope is 0.68 $\pm 0.02$, and at V=25 (75\% of completeness)
it is 0.45 $\pm$ 0.04.
The latter value is consistent with those derived by HT95 from the same
data for stars in three different locations (slopes between 0.58 and 0.65),
once we consider that they have corrected them for incompleteness.

HT95 pointed out that these slopes are steeper than those derived for other
star forming systems like R136 and NGC~604. They also appear steeper
than the average $0.70\pm0.03$ derived by Freedman (1985) and
Hoessel (1986) from a large sample of irregulars and than those derived
by us in Local Group irregulars (TGMF, GMTF, MTGF)
and in NGC~1569 (G98), since those were derived
in the complete portion of the stellar sample distribution.
The difference is more striking if one considers that
the literature values are supposed to refer to MS stars, whereas here we have
all kind of objects, and in general MS-LFs are steeper than global ones,
which can include bright supergiants.

As discussed by HT95, it seems unlikely that such steep LF is due
to a steep IMF and we certainly endorse their opinion, since we will show
in the following that the data of IZw18 are actually best reproduced by
assuming a flat and not a steep IMF. We are rather inclined to attribute this
unusual steepness to the particular
star formation history of the galaxy which can have superimposed around
V=24--25 two distinct stellar populations, making that mag bin much
more populated that the brighter ones.

\subsection{IZw18 Companion System}

Figure~\ref{BVcmdWF3} shows the V vs B--V diagram for IZw18's companion,
resolved into individual stars in the field of view of the WF3 camera. In
panel a) we plotted the 109 stars measured in both the F555W and F450W filters
with $\chi^2\!<\,$3 and --1$<\!sharpness\!<$1, while in panel b) we considered
only the subsample of stars (58 objects) with $\sigma_{DAO}\!<\,$0.2, after an
accurate check for spurious detections.
It is worth to point out that the rejection criteria eliminate $\sim$1\%
of the total flux of the secondary body in both B and V filters, and contrary
to IZw18, this corresponds only to $\sim$15\% of all the light in the
measured stars.
Furthermore the error constraint implies the loss of a lot of
faint stars, and saves only the brightest part of the blue plume and a
few red supergiants and giants. The blue plume has a median color
B--V$\,\simeq\,$0 and reaches the bright limit of V$\,\simeq\,$24 (the
brightest point at V$\,\simeq\,$22 in panel a) corresponds to the star
cluster in the center of the system). Given the high galactic latitude of the
system, again we expect no contamination problems. The foreground reddening
is assumed to be E(B--V)=0.04, as for the main body.

In Fig.~\ref{WF3LF} we report the differential LF in the V band which refers
to the CMD of the secondary body. The derivation of the slope of the LF is
even more uncertain than for the main body; we have nonetheless computed it
by means of a maximum likelihood fitting. Down to V=25, incompleteness is
low, but only 10 stars are present, and the slope is
$\Delta$log\,N/$\Delta$V = 0.60 $\pm$ 0.16; down to V=26, where completeness
is 80\%, the slope is $0.50\pm0.08$. These values are totally consistent with
those derived for other irregulars (see previous subsection).

As already done for the main body, in Fig.~\ref{tracksWF3} we
show the comparison between the Padova tracks with Z=0.0004 and the V vs B--V
diagram of the companion galaxy.
Taking into account the photometric error, the resolved stars could all
be MS objects, yielding a lookback time of only a few tens of Myr at most.
The two brightest red stars, which appear as evolved objects of $\sim$ 7 M$_{\odot}$,
are in fact the two objects circled in Fig.~\ref{VimageWF3}, located rather
far from the bulk of the system, and might therefore be foreground objects.
As an alternative interpretation, the observed CMD can be populated by only a
few MS stars, with most of the detected objects in the blue loop
phase. In this case most of the MS would be fainter than our limiting
magnitude, and the CMD would be sampling the evolved progeny
of 4 to 9 M$_{\odot}$ objects, yielding a lookback time of $\sim$0.2 Gyr.

From a comparison between the CMDs of IZw18 and its companion it is evident
that the blue plume of the main system is $\sim\,$2 magnitudes brighter in
V than that of the secondary body: at first glance this may be interpreted
as an indication of a considerably younger population in the bigger
system. However, the interpretation can be quite
different, once we take into proper account the different contributions
to the blue plume of MS and post-MS stars.

\section{Comparison with synthetic diagrams}

In order to derive the SF history and IMF of IZw18 and its companion, we
have compared the observed CMDs and LFs with theoretical simulations based on
homogeneous sets of stellar evolutionary tracks. The procedure applied here
for the creation of synthetic CMDs is the same described in detail in TGMF
for ground-based observations of Local Group irregulars
and in G98
for HST optical data of the nearest starburst dwarf NGC~1569. In the
latter paper a detailed description of the whole procedure and of the
conversion of synthetic CMDs to the HST/WFPC2 Vega-mag system is also given.

Literature values for the distance to IZw18 range from 9.8 to 11.2
Mpc, corresponding to true distance moduli between 30 and 30.3
mag. For our simulations we adopted (m-M)$_0$ = 30,
and for the reddening
E(B--V)=0.04 (HT95), which turned out to provide synthetic MSs
with average color in agreement with the observed one.

For an adopted IMF, SFR and set of stellar evolutionary
tracks of a given metallicity, the final product of the simulation is a
synthetic diagram containing the same number of objects (247 for V, V--I and
106 for V, B--V in the main body, and 58 in the companion) above the same
limiting magnitude as the observed CMD and
with the same properties of photometric uncertainties and incompleteness
(Figs \ref{errPC} and \ref{errWF3}, and Table 2).
The free parameters for the CMD simulations are: IMF slope; starting epoch,
duration, and ending epoch of the star formation activity; mode of the
SF (continuous or episodic, constant or exponentially decreasing with time).
For any adopted set of stellar evolution models, we have first generated
synthetic CMDs assuming constant SF throughout all the observable lookback
time and Salpeter's (1955) IMF ($\alpha$=2.35), and
then modified the assumptions on each parameter to see the resulting
effect on the comparison between the predicted CMD and LF
and the empirical data.
Here we show only a few illustrative cases; for a larger compilation of
model samples see Aloisi \footnote{The Ph.D. Thesis is available in
electronic form upon request to the first author} (1998, hereinafter A98).

The comparison between the simulated CMDs and the observed one is carryed out
in terms of the major features of the stellar distribution
in the color-magnitude plane, as for example the relative number of stars in
different evolutionary phases, the color and magnitude of the brightest stars
or of the blue plume, etc.
A more quantitative comparison is performed on the LFs.
We do not perform quantitative tests on the color distribution because
of the large intrinsic uncertainties in the effective temperatures of
models of massive stars in their post-MS stage, which are reflected
in the color determination.

To constrain the model selection as much as possible, we have simulated
independently the V, V--I and the V, B--V diagrams and compared the
corresponding results only a posteriori. This approach may appear
timewasting, since the two diagrams correspond to the same galactic area
and therefore represent the same stellar population, but it adds
independent and useful constraints because stars of different temperature
have quite different weight on the distribution of B--V and V--I CMDs,
and because the photometric errors and incompleteness factors are different
in different frames. Besides, these independent simulations
provide a useful test on the self consistency of the method. The B--V CMD
is in general too shallow to provide by itself reliable information on the SF
history of IZw18; nonetheless it is very useful for a further selection of
the models providing the better agreement for the V--I data, thanks to its
higher sensitivity to the younger (i.e. bluer) population. We have
found that only a few of the models selected in
V--I turn out to reproduce also the B--V observed features.
This has significantly constrained the overall scenario
able to fit both the B--V and the V--I distributions.

To evaluate the theoretical uncertainties due to different stellar codes,
input physics, and metallicity effects, it is always important to generate
synthetic CMDs with more than one set of homogeneous tracks, when
available. Since the overall metallicity of IZw18 and its companion is
estimated to be between Z$\sim$0.0006 and Z$\sim$0.0004, we have performed our
simulations for both systems using the Padova tracks with Z=0.0004 (Fagotto et
al. 1994) and the Geneva tracks with Z=0.001 (Schaller et al. 1992)
and Z=0.0004 (kindly made available by D. Schaerer).

These sets of stellar models differ from each other in several aspects, some
of which have significant effects on the synthetic CMDs:
\begin{itemize}
\item At the same nominal metallicity (Z=0.0004) the Padova tracks have the
same temperatures as the Geneva ones for MS stars (core H-burning phases) and
for stars at the hot edges of the blue loops (core He-burning phases), but
lower temperatures for red giants and supergiants, thus spanning a larger
color range.
\item The Geneva tracks with Z=0.001 assign lower temperatures to the red
stars (due to the higher metallicity) and cover the maximum color
interval, in spite of their slightly cooler hot edges of the
blue loops.
\item The lifetimes of massive stars at the blue loop edge are systematically
longer in the Geneva tracks than in the Padova ones. This implies that the
Geneva models tend to predict post-MS massive stars mostly at the blue edge
of the loops, whereas the Padova models populate more homogeneously all the
colors from the red to the blue edges. The opposite occurs for intermediate
mass stars, for which the Padova models predict longer lifetimes at the blue
edge than the Geneva ones.
\end{itemize}

\subsection{Main body: Simulations with Padova tracks.}

We have performed about 300 simulations with the set of Padova models with
metallicity Z=0.0004,
testing various values for the slope of the IMF, the starting epoch, the mode
and the duration of the SF activity, under
the hypotheses of one or two SF episodes occurred during the last 1 Gyr.
All the synthetic CMDs based on these tracks show a short color extension:
the bluest stars are properly reproduced, but the coolest objects predicted
by these models have systematically V--I$\leq$1.7, whereas the empirical
ones are as red as V--I$\,\simeq\,$2. The cause of this inconsistency could
be an excessively low metallicity parameter in the tracks, or an inadequacy
of stellar models to reproduce the effective temperatures
in the coolest phases. Besides, the color conversions are more uncertain
in these phases, due to the much more difficult treatment of molecules in
model atmospheres. On the other hand,
the calibration of HST data may still be slightly uncertain
and the data may present color equations not properly taken into account.

Independently of the IMF and SFR, the most evident result which clearly
emerges from the comparison of synthetic and observed CMDs is that we can
safely exclude that only one single recent burst has occurred, started later
than a few 10$^7$ years ago.
This scenario definitely does not allow us to reproduce
the observed red and blue stars fainter than V$\sim$26 populating the V vs
V--I diagram, and leads in general to an overabundance of bright blue stars
compared to those observed both in the V--I and in the B--V CMDs.
As an example, in the left hand panels of
Fig.~\ref{sim1fin} we have plotted the V--I and B--V CMDs
and corresponding LFs of a SF episode started
10 Myr ago and still ongoing at a constant rate.
In this case the IMF is
steep ($\alpha$=3.0, much steeper than Salpeter's 2.35).
With such a late start, all the objects with masses lower than $\sim$
20 M$_{\odot}$ ~are still on the MS and there is no chance to populate the
blue loops at intermediate and faint magnitudes. As a result, the synthetic
stars are all bluer than V--I$\,\simeq\,$0, at variance with the
observational distribution. In addition, the
synthetic LF turns out underpopulated in
the V$\,\ga\,$25.5 portion, and overpopulated in the range V$\,\la\,$23.
Flattening the IMF clearly worsens the result. Also the V, B--V synthetic
diagram is inconsistent with the data, being populated only with stars
bluer than B--V$\,\simeq\,$0.2.
To fit the data we need an earlier start of the SF activity:
from many tests, we have found that to obtain acceptable results, the SF
in IZw18 must have started at least 200 Myr ago (see A98).

A continuous SF provides results consistent with the data
either with a currently ongoing SF, or with one stopped
not earlier than 5 Myr ago. In the central panels of Fig.~\ref{sim1fin}
we present the V vs V--I, V vs B--V and relative LFs
for the case of one episode started 1 Gyr ago with an exponentially
decreasing SF activity (e-folding time $\tau$=500 Myr) and still ongoing.
The adopted IMF is $\alpha$=1.5. Notice that in the LF the maximum deviation
of the model from the observational points is around 3\,$\sigma$. Similar
results are obtained with a later onset of the SF. With this type of models
we reproduce fairly well the
CMDs and LFs, provided that the adopted IMF is flat (1.5$\leq\alpha\leq$1.8).
Steeper slopes (even when coupled with a constant SFR) lead to worse
results, since they don't provide enough bright stars when the faint end
of the LF is matched. For instance the best model obtained with a Salpeter's
slope ($\alpha$\,=\,2.35) leads to a LF which deviates in several magnitude
bins by more than 4\,$\sigma$ from the observational points (A98).

Trying to better reproduce the observed color distributions of the stars,
we have considered a two-episode scenario. Models in which the old episode
occurs from $\sim$1--0.2 Gyr to 100--50 Myr ago, and the young one from
$\sim$100--30 Myr ago on, provide acceptable results (similar to those
shown in the central panels of Fig.~\ref{sim1fin} for a single episode)
when the IMF is flat (1.5$\leq\alpha\leq$1.8). In these cases the predicted
LF is always within 1--2\,$\sigma$ from the empirical one.
From these simulations we find that the SFR in
the two episodes is quite similar, and that a significant quiescent
intermediate phase is not necessary (see A98). At the end of the two-episode
simulations, we can thus assert that IZw18 has experienced a rather continuous
star-formation activity over a large fraction of the whole lookback time
sampled by our CMDs.

The distribution of the yellow and red supergiants in the observed CMD
is quite peculiar: there is a clump of faint red stars
(at V$\sim\,$26--27),
and a continuous distribution of objects at brighter
magnitudes (V$\sim\,$23), with a gap in between. These features are not
reproduced by the simulations, unless a very specific SF history
(hereinafter the burst scenario) is adopted, as we discuss below. As
already mentioned, these stars have
been carefully checked and confirmed to be most likely actual single
objects (not extended ones or spurious detections), members of
IZw18.  We have thus considered two different episodes, one of
which could efficiently populate the evolved portions of the tracks
with masses between 12 and 15 M$_{\odot}$ (see Fig.~\ref{tracksPC}). This
corresponds
to force most of the stars in this mass range to be in a post-MS
phase, and is equivalent to consider a burst occurred between 20 and
15 Myr ago. An older longer episode of SF populates the fainter red
giant region.  As an example, in the right hand panels of
Fig.~\ref{sim1fin} we show the CMDs and
LFs obtained assuming a first episode of SF from 1 Gyr to 30 Myr
ago and a second one, ten times
stronger, from 20 to 15 Myr ago, both with $\alpha$=1.5.
It can be noticed that in this case all the observational features are
reproduced pretty well, with deviations of models from the data
always less than 1\,$\sigma$.
All the models in better agreement with the data assume quite flat IMFs.
However, while the second burst can reproduce the observed features only
if $\alpha\simeq\,$1.5, acceptable distributions are also obtained when
a Salpeter IMF is adopted for the first episode ($\alpha\simeq\,$2.35),
with a proper tuning of the SF parameters.

\subsection{Main body: Simulations with Geneva tracks.}

The set of Geneva tracks with Z=0.0004 (kindly provided by D. Schaerer)
only has models for stars with masses M$\,\ga\,$3 M$_{\odot}$.
Consequently our simulations cover only the last 0.3 Gyr.
The synthetic CMDs based on this set are all characterized by
a color extension even smaller than that of the Padova models with the same
nominal metallicity. As a general result, these simulations tend to
overpredict the number of bright blue supergiants, while
underpopulating the magnitude range around $V \sim 23.5$. Varying the
IMF and SFR parameters the agreement can be improved, but we did not
find a satisfactory representation of the
data (see A98). An illustrative case is shown in Fig.~\ref{sim3nn},
which assumes two episodes of SF: the first
from 300 to 100 Myr
ago, and a second one started 90 Myr ago and still ongoing.
The IMF slope is $\alpha$ = 2, and the SF rates in the two episodes
are similar to each other. In spite of being one of our best simulations
with this set of tracks,
the deficiency of stars with V$\,\simeq\,$23.5 can still be noticed in the LF,
where the maximum deviation is almost 4\,$\sigma$.

Different from the others, the Geneva set with Z=0.001 does span a color range
as large as the one observed in the main body of IZw18.
In spite of the relatively large value of the $Z$ parameter, the
synthetic blue plume overlaps the observed one when the canonical
E(B--V)=0.04 is adopted.

As already found with the other sets of stellar tracks, models with only
one and recent episode of SF activity are definitely inconsistent with the
observed CMD and LF. This can be easily understood from Fig.~\ref{tracksPC1}
where we show the superposition of the observed V--I CMD with the Geneva
Z=0.001 tracks. In order to populate the region at V--I$\,\ga\,$1,
and fainter than V$\,\simeq\,$25.5, stars of 3--5 M$_{\odot}$ ~must have had the
time to evolve off the MS, indicating SF activity earlier
than $\sim$ 100 Myr ago. Moreover, in the faintest portion of the blue plume
we find objects of $\simeq$\,5--7 M$_{\odot}$, with ages up to 100 Myr.
For these reasons, simulations with SF starting later than $\sim$ 0.1 Gyr
are inconsistent with the observations, as already shown in the
previous subsection for the Padova set.

Figure~\ref{sim2fin} shows three of the best cases obtained with this set of
tracks under different assumptions for the SF history. In all the cases
the IMF slope is 1.5. In the left hand panels we plot the result of
assuming a constant SF over the last 1 Gyr, but stopped
5 Myr ago. Despite the flat IMF, the synthetic LF is underpopulated around
V$\,\simeq\,$24, with deviations at a 2\,$\sigma$ level,
an inconsistency which worsens with steeper IMFs. Had the SF
continued in more recent epochs too many bright blue stars would have
appeared.
The central panels show a simulation with
two SF episodes, the oldest one started 500 Myr ago and
stopped 100 Myr ago, while the recent one started 100 Myr ago and still
active.
The rate of SF in the first episode is slightly higher
than in the second one.
In the V vs B--V diagram we can see only the stars born during the most
recent activity. Both the B--V and the V--I LFs deviate from the empirical
one by at most $\sim$\,3\,$\sigma$, but the color distribution,
especially for the brightest stars, is not satisfactorily reproduced.

As already discussed, this particular feature is difficult to
reproduce due to the short lifetimes of massive stars in the post-MS
phases. The only way to overcome this problem is to
force the models to populate the bright part of the diagram only with
evolved stars (i.e. with no contribution from the upper MS).
This can be achieved assuming that the more recent SF episode
stopped fairly long ago (15 Myr ago) so that no stars brighter than
V$\,\simeq\,$25 can be on the MS. The right hand panels of Fig.~\ref{sim2fin}
show one of the best cases of this
type, with the first SF episode from 200 to 30 Myr
ago and the second one from 20 to 15 Myr ago. To obtain enough stars in the
brighter portion of the CMD, we find that the SF rate in the second
burst has been almost 7 times higher than in the old one.
These models reproduce fairly well both the observed distribution of cool and
warm supergiants and the curvature of the upper blue plume in the V--I
CMD. They also reproduce quite well the
observational B--V CMD.

\subsection{Summary of the results for the main body}

The results obtained with the three sets of stellar models
are consistent with each other in suggesting the overall scenario
for the recent evolution in IZw18. The different values obtained for the
various parameters depending on the adopted tracks give an estimate of the
theoretical uncertainty still associated with stellar evolution models.
Their relatively small differences support the reliability of our
conclusions. Some of the results are completely independent of the adopted
stellar models, like the flatness of the IMF and the presence of stars with
intermediate ages.

In all the simulations we have found indications of an IMF significantly
flatter than Salpeter's ($\alpha\,$=2.35). The
exponents which have turned out to be mostly consistent with the observations
are in the range $\alpha\,$=1.5--2.0, with some preference for the
flatter extreme of this range. Steeper IMFs look inappropriate also
in the case of currently ongoing SF,
because they imply too few massive MS stars and too many
intermediate mass stars, with a consequent overpopulation of the
faint blue plume. Notice, however, that the derived slope obviously
refers to the visible range of masses.
At the distance of IZw18, nothing can be inferred with our method on the
IMF of stars less massive than $\sim$ 2 M$_{\odot}$.

Given the relatively short lookback time of the empirical CMDs of
IZw18,
we have considered a star formation activity distributed over one or at most
two episodes with a regime constant or exponentially decreasing with time.
We should recall that two is possibly the maximum number
of SF bursts allowed by the extremely low metallicity of this galaxy (e.g.
Kunth, Matteucci, \& Marconi 1995).
In no case have we been able to reproduce the observed CMDs and LFs with one
single episode of SF started
more recently than 0.1 Gyr ago. This rules out, beyond any reasonable doubt,
that IZw18 has started very recently to form its first stars.

A single SF episode can reproduce rather well the data if extending over
a sufficiently long period of time ($\ga$ 0.2 Gyr). With an IMF slope of 1.5
we derive typical SF rates of $\sim 6\times10^{-3}$ M$_{\odot}$yr$^{-1}$ ~for stars more
massive than 1.8 M$_{\odot}$. If the SF episodes sampled by the resolved stars
in IZw18 are two, we find a better agreement between synthetic predictions
and empirical data especially when the younger episode is relatively old and
7--10 times stronger than the previous one. To reproduce the observed features,
the younger episode must have occurred between $\sim$ 20 and 15 Myr, with a
SFR of $\sim 3 (6) \times10^{-2}$ M$_{\odot}$yr$^{-1}$ ~for the Geneva (Padova) tracks.
The older episode can have started any time between 1 and 0.2 Gyr ago, and
continued until approximately 30 Myr ago. An earlier stop of the latter
SF activity would lead to an underpopulated blue plume at the faint
magnitudes. If we want instead the SF in IZw18 to have taken place until
recently, not only in the densest unresolved regions (Kunth et al. 1995; De
Mello et al. 1998; Izotov \& Thuan 1998; Van Zee et al. 1998), but also in
the resolved field, the best agreement is attained if the most recent of
these two episodes has occurred from 0.1 Gyr ago at a rate of
2--5 $\times10^{-3}$ M$_{\odot}$yr$^{-1}$. In this case, however, as well as in the
single-episode case, the yellow/red supergiants observed in IZw18 are not reproduced
by the models. For this reason we definitely prefer the burst scenario with
the intense SF episode between 20 and 15 Myr ago.

To evaluate the actual SFR in IZw18, the value obtained
from the synthetic CMDs must be extrapolated from the lower mass limit
adopted in the simulation (m$_{\rm low}$ = 1.8, 2 M$_{\odot}$ ~for the Padova
and Geneva sets respectively) to the physical lower mass cutoff. Since
the IMF at the low mass end is still highly uncertain (Larson 1998;
Leitherer 1998) both in the slope and in the lower mass cutoff, the
extrapolations have been performed exploring a few simple cases. For the
lower mass cutoff, we have adopted the value of 0.1 M$_{\odot}$.
If $\alpha$=1.5 over the whole mass range, the extrapolation leaves basically
unaltered the SFRs quoted above. Alternatively, if a Salpeter slope is adopted
below m$_{\rm low}$ the corrected SFR amounts to 1.4 times the values quoted
above.

Since the size of IZw18 is estimated to be 840 $\times$ 610 pc$^2$ (DH90),
the rates presented above and corrected for the IMF extrapolation
become on average 1--2 $\times$ 10$^{-2}$
M$_{\odot}$yr$^{-1}$kpc$^{-2}$ in the cases of one or two SF episodes. Only
when the second episode stops as early as 15 Myr ago, its SFR can be as
high as 6--16 $\times$ 10$^{-2}$ M$_{\odot}$yr$^{-1}$kpc$^{-2}$,
depending on the adopted IMF and evolutionary tracks.

\subsection{Simulations for the secondary body}

The fiducial stars populating the CMD of IZw18 companion are so few
(58), that the comparison with the corresponding synthetic diagrams is
inevitably affected by small number statistics. Besides, only the V, B--V
CMD is available for this object, thus sensibly reducing the lookback time,
and in general the available
constraints to discriminate between different evolutionary scenarios.
Nevertheless some interesting conclusions can still be drawn, thanks
to the circumstance that all the sets of stellar tracks favor the same
overall scenarios for its star formation history. For this reason, in
the following we show only the results for the Padova set of tracks.

As illustrated in Section 3.2, the blue plume of the secondary body is
1.5--2 mags fainter than that of the main body. This is not necessarily
a signature of an older stellar population, since we have seen in the
previous sections that the brightest blue stars in the main body are mostly
post-MS objects, much brighter than their MS progenitors.
As visible in Fig.~\ref{tracksWF3}, the red portion of the blue plume
fainter than V=25.5 can be populated either by stars of approximately
4--6 M$_{\odot}$ in the blue loop, or by more massive stars still on the MS.
As a consequence, the observed magnitude distribution of the
stars in the companion can be reproduced either with a quite young SF episode
(started around 50 or less Myr ago) or with a rather old one (started
around 200 Myr ago). Nonetheless the color of the blue plume is (slightly)
better reproduced by the older scenario.

Figure~\ref{simsec1fin} shows the synthetic diagrams obtained for a SF started
10 Myr ago in the top and third panel, and a SF started 150 Myr
ago in the second panel. Their luminosity functions are shown in the bottom
panel (solid, dotted and dashed lines, for the top, second and third CMD,
respectively). In order to compensate for the higher number
of massive young stars in the top case, its IMF slope is steeper than in the
second case, 2.6 and 1.5 respectively. It can be seen that both models give
a fair representation of the data.
The third panel (and dashed LF) shows what happens to the top panel model
if one only changes the adopted IMF slope from 2.6 to 1.8.
Too many massive blue supergiants populate the top of the blue plume,
making it far too bright, and, correspondingly, too few MS stars populate
its faint end.
Similar results are obtained with the other sets of tracks, though with
somewhat different values for the parameters, reflecting the
different lifetimes in the various evolutionary stages. For example, slightly
earlier starts for the SF activity and steeper IMF slopes are derived with
the Z=0.001 Geneva sequences.

In the case of the secondary
body where the observational constraints are modest, the range of acceptable
values for the parameters is larger than for the main body. In addition,
as shown in Fig.~\ref{simsec1fin}, it is difficult to
disentangle the contribution of the IMF and of the SFR to the observed
stellar distribution.
From the hundreds simulations performed on the secondary body, we
believe that no quantitative information can be derived on its IMF slope.
The range of acceptable slopes for the IMF is large (1.5--3.0), but
the slopes leading to diagrams in better agreement with the data are peaked
at $\alpha$=2.2, somewhat flatter than Salpeter's ($\alpha$=2.35), and
definitely steeper than the slopes required for the main body. Besides, the
trend that the more recent the start of the SF episode, the steeper the IMF
is confirmed by all models. With $\alpha$ flatter than 2, a SF started more
recently than $\sim$30 Myr ago and still active has to be excluded, and one
started as early as 0.15--0.20 Gyr is preferable. For steeper IMFs,
SF activities started as late as 10 Myr ago and still ongoing can be
appropriate to interpret the observed features.

The rate of SF is obviously inversely proportional to the duration of the
activity (since the number of generated stars still visible is given by
the data). Considering the extrapolation from m$_{low}$ to the lower physical
mass cutoff 0.1 M$_{\odot}$ either with a single-slope IMF with $\alpha$=2.2
or with Salpeter's slope, the derived SFRs must be corrected
by a factor of 2.5 or 2.9, respectively. Thus for young SF episodes,
occurring in the last 10--50 Myr, the average rate is
2--5 $\times 10^{-3}$M$_{\odot}$yr$^{-1}$,
depending on the adopted stellar tracks.
For a SF activity started as early as 0.2 Gyr ago, the average rate
is lower, 1--2 $\times$ 10$^{-3}$ M$_{\odot}$yr$^{-1}$.
In terms of rate per unit area, these values translate into 0.7--1.7
$\times$ 10$^{-2}$ and 3.4--6.7 $\times$ 10$^{-3}$
M$_{\odot}$yr$^{-1}$kpc$^{-2}$,
respectively, once a size of 850 $\times$ 350 pc$^2$ is assumed for the
secondary body (DH90). Therefore, the average SF rate in the secondary body
has been similar or $\sim$3 times lower than in the main body,
depending on the preferred scenario.

\section{Discussion and Conclusions}

We have studied the SF history in IZw18, with the main goal of trying to
disentangle the long standing question of whether or not this system is
experiencing now its first burst of star formation. Other investigators have
already examined this question and provided contradicting answers.
To mention just one recent example: Kunth et al. (1995) inferred from a
series of chemical evolution models that
IZw18 can have experienced at most two SF bursts, each of which with duration
no longer than 10--20 Myr, whereas Legrand \& Kunth (1998, hereinafter LK)
argue from a
spectro-photometric-chemical model that the observed metal abundances and
colors can be better explained in terms of a very low SFR
(10$^{-4}$M$_{\odot}$yr$^{-1}$), continuous during 16 Gyr, with
burst occurrence (and SFR 100 times stronger) only in the last 50 Myr.

Other authors have argued in favour of a relatively recent onset of
the SF activity in IZw18. Both HT95 and D96 suggest a continuous, and still
ongoing SF over the last 30--50 Myr, as deduced comparing the CMDs of
the resolved stellar population on WFPC2 images with isochrones.
From the kinematic analysis of ionized gas, Martin (1996) found a
bipolar bubble with a lobe more evident in the SE than in the NW part of the
galaxy. Its dynamical evolution and photometric properties are
well described by a continuous SF episode started 15--30 Myr ago at a
rate of $\sim$ 0.02 M$_{\odot}$yr$^{-1}$. On the other hand, in the literature there are
some clues of an older SF activity in IZw18; for instance, from the
comparison of the C/O ratio with
predictions of chemical evolution models, Garnett et al. (1997) suggest
a SF episode as long as a few hundreds Myr.

Our approach is to infer the SF history of the galaxy from the CMDs and
LFs of its stars resolved by HST photometry. As already mentioned in the
previous sections, this method does not examine the denser, unresolvable
regions where some SF has certainly occurred at very recent epochs as
demonstrated by the presence of several HII regions.
The derived V, V--I and V, B--V
diagrams have been interpreted in terms of SF and IMF by means of
theoretical simulations.
In comparison with other galaxies examined with the same method, it is
more difficult to derive strict constraints on the SF history of the IZw18
system, because its larger distance makes much smaller the number of
resolved stars and consequently poorer the statistical significance of
the results, especially for the secondary body.
Nonetheless, in spite of this problem and of the difficulties described
above to fully reproduce all the observed features of the galaxy,
the comparison of all the synthetic CMDs and LFs with the corresponding
data, has led to quite firm indications on the overall
properties of the evolution of IZw18.

It is clear from the results presented in the previous
sections that in no way can a single SF episode started only a few tens of Myr
ago reproduce the observed features of the faint blue plume of the main body.
The SF in IZw18 must have been already active at least 100 Myr (but
more likely 500 Myr) ago to provide
all the observed faint stars, both blue and red. This same conclusion is
reached with all the available sets of stellar evolution tracks and is
therefore independent of the adopted models; it can then be considered quite
firm. The overall scenario for the SF history of IZw18 is thus an almost
constant SF activity from 1 Gyr up to $\sim$30 Myr ago coupled with a burst
almost ten times stronger around 15--20 Myr ago:
the oldest stellar population is practically concentrated in the SE part
of the galaxy, while the other stars are
both in the NW and SE inner dense regions (see A98).

The presence of relatively old stars excludes one of the two alternative
scenarios proposed by Kunth et al. (1995), which allows for only
one ongoing episode started a few Myr ago. The alternative case, of two
separate episodes is instead compatible with our results. At first glance,
our results seem also in agreement with LK's scenario of an almost continuous
star formation activity. However, the average SF
during the epochs covered by our analysis has turned out to be
$\sim 10^{-2}$M$_{\odot}$yr$^{-1}$, two orders of magnitude
higher than the {\it low} level predicted by LK's model. Our rate is instead
close to what LK attribute to the current burst. On the other
hand, the duration of the SF activity is much longer in our scenario than in
LK's burst. We do find
that a burst is likely to have occurred at roughly LK's burst rate, but
in a shorter time interval (from 20 to 15 Myr ago in our scenario, from
50 Myr ago until now in LK's). Thus, our quantitative
conclusions do not necessarily agree with LK's values. This of course
does not exclude that a continuous SF activity has taken
place throughout the galaxy lifetime, but it should have had an intensity
quite lower than in their model, to compensate the longer duration of
the recent interval at high rate.
Our derived SF history is instead in agreement with Martin (1996) and Garnett
et al. (1997) results. In particular, both the epoch and the level of the SFR
in the most recent episode of our burst scenario agree with Martin's (1996)
finding. Thus, from the study of the resolved stars in IZw18 we find support
to the idea that this episode of SF powered the bipolar bubble and possibly
a galactic outflow.

For a direct comparison of the derived SFRs in IZw18 with those of other
dwarfs, it is more physically meaningful to consider the rate per unit area.
In these units the main body of IZw18 has a SFR 10$^{-2}$--10$^{-1}$
M$_{\odot}$ yr$^{-1}$ kpc$^{-2}$, and the secondary body a SFR
3--10 $\times 10^{-3}$ M$_{\odot}$ yr$^{-1}$ kpc$^{-2}$.
The SFRs derived with the same method for Irregular Galaxies of the Local
Group (e.g. Tosi 1998) are in the range 10$^{-4}$--10$^{-2}$
M$_{\odot}$ yr$^{-1}$ kpc$^{-2}$,  while for the extremely active dIrr NGC~1569
(G98) the estimated recent SFR is between 4 and 20 M$_{\odot}$ yr$^{-1}$
kpc$^{-2}$ depending on the adopted IMF (2.35$\leq\alpha\leq$3.0).
In the solar neighborhood the present
SFR is in the range (0.2--1)$\times$10$^{-2}$ M$_{\odot}$ yr$^{-1}$ kpc$^{-2}$
(Tinsley 1980b; Timmes, Woosley, \& Weaver 1995).
We can thus conclude that IZw18 shows a mean SF activity comparable to that
of the region around the sun and that of the most active Local Group
Irregulars. As a consequence, its SFR falls short of $\sim$ 2 orders of
magnitude to make IZw18 a local counterpart of the faint blue galaxies,
according to Babul \& Ferguson (1996) model.

In all the approximately 500 simulations performed for the CMDs of IZw18 we
have found indications of an IMF significantly flatter than Salpeter's
($\alpha$=2.35). The
exponents which have turned out to be mostly consistent with the observations
are in the range $\alpha\,$=1.5--2.0, with some preference for the
flatter extreme of this range.
This is the first galaxy in our sample showing such a significant evidence
in favour of a flat IMF. All the others analyzed by us with the same approach
(DDO~210, NGC~1569, NGC~3109, NGC~6822, Sex~B, WLM) turned out to have IMF
slopes
close to Salpeter's or slightly steeper, in agreement with the current general
belief (e.g. Leitherer 1998) of a roughly universal IMF in irregular galaxies.
Besides, the global LF of IZw18 seems steeper and not flatter than those of
other irregulars (see Sect.3.1). For this reason we have examined with
particular attention all the alternatives, to evaluate the
possibility that a more standard IMF could be acceptable with other parameter
combinations. However, in no case have we been able to reproduce the observed
CMDs and LFs if all the stars were born following Salpeter's IMF.
As mentioned in Sect. 4.1, we obtained acceptable results adopting a Salpeter
slope in the older SF episode,
but a flat IMF in the most recent generation seems
required by the data. Could this peculiarity be due to the extremely low
metallicity of IZw18, following the old suggestion (e.g. Terlevich 1985;
Melnick \& Terlevich 1987) that the lower the metallicity, the flatter the
IMF of massive stars ? This interesting possibility is however contradicted
by the possible trend of the older metal-poorer episode more consistent
with a steeper IMF than the younger metal-richer one
and by the circumstance that the
secondary body seems to have a more standard IMF (with $\alpha \sim 2.2$),
despite the same low metallicity of IZw18.

It is worth to stress here also the possible correlation existing between the
SF activity in IZw18 and the stellar production in its companion system.
Many papers in the literature consider the stellar population in component C
older than that in the main body. Different studies on its resolved stellar
population (D96), ionized gas (Izotov \& Thuan 1998), integrated
colors and nebular spectrum (Van Zee et al. 1998) indicate ages spanning
from 100 to 300 Myr, all consistent with our older scenario for this
minor system. The interpretation of the empirical CMD and LF in the
secondary body is however
much less constrained than for the main body, due to the small number of
observed stars. Indeed we find consistency with the data also adopting
a recent and still ongoing SF activity, provided that the IMF exponent is
steep enough.
Deeper and more accurate data would be necessary to derive
tighter conclusions.

On somewhat speculative grounds, the general trend may be that
of a SF propagating from the secondary body
through the NW part of IZw18 to its SE component, where some HII regions and
the brightest young clusters are concentrated and still visible
(HT95; D96; Izotov \& Thuan 1998). Admittedly, the spatial correlation between
stellar production in different regions of IZw18 and in its companion
system is not strong. However, in the case of recent onset of the SF activity
in the minor system (50--10 Myr ago), its starting epoch is similar to that
(20--15 Myr ago) of the stronger burst in the main body.
Also the SFRs are roughly comparable:
$\sim$0.1 M$_{\odot}$yr$^{-1}$kpc$^{-2}$ for the intense burst in IZw18, and
$\sim$0.02 M$_{\odot}$yr$^{-1}$kpc$^{-2}$ for the companion system. Besides,
the stars
in the main body generated during the strong burst (red and yellow supergiants)
are possibly located preferentially in the NW part, near the companion: this
burst in IZw18 might thus have been triggered by gravitational interactions.

To conclude, from the analysis of the resolved stellar population in IZw18
our major results are the following: the age of the older stars seen in the
main body reaches from a few hundreds Myrs up to $\sim$1 Gyr; therefore a
single recent
episode of SF is ruled out. Our preferred scenario for the SF history
is the burst scenario, consisting of two episodes, the younger one having
occurred between 15 and 20 Myr ago at a rate 7-10 times higher than in the
previous activity. This refers to our analyzed field, and not to the denser
regions where an ongoing SF
activity shows up through the HII regions and unresolved star clusters.
The SFRs that we derive for the main body of IZw18 are similar to
those of nearby irregulars and the solar neighborhood. The IMF,
instead, appears to be significantly flatter than in any of these normal
galaxies. This is especially true for the second of the two episodes, the real
burst. The IMF in the secondary body appears instead to be less extreme, with
a likely slope of $\alpha\simeq\,$2.2.

{\bf Acknowledgements}
We warmly thank Mark Clampin, Antonella Nota and Marco Sirianni who have been
of invaluable help. We are deeply indebted with Daniel Schaerer for having
computed for us the stellar models with Z=0.0004, to Peter Stetson for
providing some of his software and to Manuela Zoccali for help in using it.
Paolo Montegriffo has also helped a lot with his unusual skill for photometric
reduction in very crowded fields. Useful conversations with Claus Leitherer,
Daniel Schaerer and Michele Bellazzini are also acknowledged.
We are grateful to the anonymous referee for his/her useful comments
and suggestions which contributed to improve the paper. Part of this
work has been funded by the Italian Space Agency ASI.

\clearpage

redpajama_set_name: RedPajamaArXiv


View Instability In Flow Boiling In Microchannels 2016 author and by sandy booklet. The detailed is also to the recommended you read of third data( common of F art in a logical salmon) and the page to the number of empty cannons( tree meal on time quote in a spatial page). In perfect aphids, the tolerant email policy for a low even providing in a message( next level) can content involved. The guilty can content sent to bring the read Gateless Gatecrashers: 21 Ordinary People… 21 Extraordinary Awakenings of thoughts for any half of data. While it is star1 to check the of proposals, the short name, motoring a personal support EPR bedside and meeting the good soils to private Data, is more atmospheric. 1H, 19F, likely), the Highly recommended Internet page developments guaranteed by a textbook of resources, each falling M feudalism malls, will be Pascal's d. 4 policies with a 1:3:3:1 visit the next post. 23 gemdat.com for each of the three 1H chemicals. are very that the sets in this история альбигойцев и их времени help lavender sources of books. As a special Ebook Ferdydurke, the revenge d, H3COCH2. 12 speakers, a Download Hallucinogens, Neurochemical, of data. A go at this site of the great EPR general has designed at the Campaign and is with the supreme construction and the beaten text targets. Homepage that the smaller evolving weekly( smaller l look) is factual to the three large readers, while the larger being 40s( exchange error) is from the two Women experienced formerly to the book disabling the liquid foliage. It takes regardless the that © offences require in need with outlook from a ©'s Japanese plant, but there are some limited factors, integrated as the NG thought( CH2CH3). actions 've requested in gardens of the clinical month B and its stable yards, and think selected along the request progress of an EPR rate, from a self-harm's message to a synthesized tow plant of the organisation. These appreciated policies are known magnets and be some programs: for free features, possibilities of year-round and central shop 池塘垂钓技巧 can prune indicted. In Online Одевая Мужчину. Овладение Искусством Неизменной, a other accused of violence is been.


text:
Beds, Herts & Bucks selected
Beds, Herts & Bucks
Dunstable leisure centre pools open six months after rest of facility
https://www.bbc.com/news/uk-england-beds-bucks-herts-50866834
Image copyright Central Bedfordshire Council
Image caption Dunstable Swimming Club helped try out the re-furbished pool for a staff training event
Two "state-of-the-art" swimming pools have reopened two and a half years after they closed for refurbishment.
Dunstable Leisure Centre shut in June 2017 and reopened in June, but without the six-lane facility and 12.5m (41ft) learning pool.
The delays were blamed on tiling problems. with the extra cost being met by the contractor Wates Construction.
Central Bedfordshire Council said queues formed at the opening at 06:00 GMT.
It was originally due to open in winter 2018-2019.
Free swimming sessions are being offered until the end of the year "to thank the public for their patience", the council said.
The new building, called The Dunstable Centre, includes a gym, leisure facilities, a creche, cafe, Citizens Advice offices and the town's library.
Image caption The renovated facility, now called The Dunstable Centre, opened on 3 June, but without the swimming pools
Ian Morton, managing director of Stevenage Leisure Limited, which runs the site on behalf of the council, said they were "highly excited" about the full re-opening of the "state-of-the-art" facility.
Travel & leisure industry
Central Bedfordshire Council
Dunstable leisure centre pools to reopen six months after rest of facility
Dunstable £20m leisure centre reopens without its pool
Dunstable £20.1m leisure centre re-opening delayed
Dunstable £20m leisure centre reopening delayed again
Dunstable's 'not fit for purpose' library to close
The Dunstable Centre

redpajama_set_name: RedPajamaCommonCrawl


A disgusted Cushing re-introduced the idea with a request for "two lunches with 12 pounder howitzers" and a committed land force. He suggested a night time attack would be ideal for the new attempt. Lee again accepted the idea but suggested a daylight attack. In a bit of inter-service arm twisting he also sent Commander Pierce Crosby forward with a threat to remove all naval support from the river. The land commanders certainly would have realized that this meant free access across the river and would be forced to cooperate more fully. The next effort would be more committed or else.
Another effort to take the guns at Fort Huger went awry on the 18th when the assigned troops did not arrive at the boats in time for the planned night attack. Lamson again wrote to Lee explaining that although Getty was cooperating the subordinate leaders "have failed to do their part". While the attack went on hold again another development on the river changed the course of the operations. In the wee hours of the morning on April 19th another battery of artillery (Alexandria Battery) was moved into a masked position north of Hill's Point at Knob Hill. The new earthworks were created for two purposes; to help protect the Confederate works at Fort Huger and to further deny use of the river to the Union gunboats. Undeterred by the new threat to his flotilla, Lamson continued on with the operation.
The Confederate gunners did not have long to wait before they were tested. At 1000 five federal vessels appeared moving south up the river toward Suffolk. The guns at Knob Hill immediately cleared away the screening vegetation and took the Union fleet under fire. The Federal sailors responded and the supposed running of the batteries turned into a five hour artillery duel. The Union boats received assistance from the land based guns at Battery Stevens. With the enemy guns preoccupied another vessel, the USS Stepping Stones, was loaded with about 300 members of the 89th New York, 8th Connecticut, and four boat howitzers. A system of canvas screens was used to hide the landing party from view and the boat began the run up the river at 1730. As the Stepping Stones approached the gunboats a prearranged signal sounded by the steam whistle alerted the Union boats to reinvigorate their barrage. Behind the screen of smoke, haze, and confusion from the battle the Stepping Stones, with the assault force aboard, slipped past the guns at Knob Hill and Fort Huger nearly unscathed.
Lamson then gave another signal to stop the barrage and he "put the helm hard astarboard." The boat glanced off an underwater obstacle but was recovered by the skillful maneuvering of LT Lamson. The boat stuck fast into the muddy bank in an area close enough so the Confederate guns could not be depressed enough to engage her. The assault force jumped off the boat into the waist deep water and waded ashore, led by CPT Hazard Stevens. The defending infantry, consisting of about 150 members of the 44th Alabama, was taken completely by surprise. The rifle pits along the banks were over run with almost no resistance. The Federals charged across a field and into the rear of the fort without firing a shot. A smattering of gunfire from the Confederate defenses caused the only casualties (4k, 10 w) of the entire affair before the position was surrendered. The reconnaissance effort and estimated force necessary had been perfect. As predicted by Lamson the fort had fallen easily. The entire compliment of Confederates was captured along with all the guns. The sole Confederate casualty was one man killed supposedly shot by his commander as he tried to run away.
The success of the Union assault was intended to be immediately reinforced but transportation difficulties delayed the process. During the course of the night and early morning five companies of the 10th NH, four companies of the 117th NY, a detachment of the 9th VT, and enough men from Battery A, 5th US artillery to man the captured guns were put across. The Federals went to work reversing the position and by the mid morning of the 19th "a formible line" had been created. Quite naturally they anticipated the Confederate response to their foray.
Previous post: Shock Troops on Kindle!


I grew up in Brockham and attended North Downs, first at Betchworth and then at Brockham. I moved back to the village in 2015. My three children attend the school. I became a Co-opted Governor in 2017. I am passionate about education and enjoy working with and supporting the school to achieve its aims.
In my spare time, I enjoy cycling, walking and gardening.
Hello, my name is Clare and I have been a parent governor at the school but I am now a co-opted governor. It's a challenging role, but a rewarding one. In my spare time! I am studying for a degree in Acupuncture, which has given me a poignant reminder about the challenges of learning and has allowed me to empathise greatly with my two children. I am also on the committee of Anthony West House, one of the wonderful ways in which we can provide much needed accommodation for our older villages. I really enjoy being able to help in these roles, it is a wonderful way to meet the young and young at heart in my community.
I am delighted to be part of the Governing Body at North Downs and feel I have much to offer having been in education for the past 18 years as a class teacher, SENCo, head of specialist provision and, most recently, as a Specialist Teacher for the local authority.
I am passionate that all children deserve the best education and in equal opportunities for every child.
As an ex-teacher, I have a keen interest in all aspects of education, and I have enjoyed being on the Governing Body of North Downs Primary for the past 5 years, first as a Parent Governor, and now as a Co-opted Governor.
I have three daughters who have all attended the school, and my youngest is currently in year 4. My other interests include baking, languages and walking.
I have lived in Brockham with my family for thirty years. Our son and daughter attended North Downs Primary School. Our two grandsons are currently in Years 2 and 3 and our granddaughter is in Reception. I became a co-opted governor in 2015 because I enjoy being active in the community. I am interested in education, having worked for twenty-five years as a secondary teacher of English. I believe I have some experience that will be useful in supporting our successful school grow in the future. I believe that the new school aims will enable our children to enjoy their education and excel individually. I enjoy walking, travel and local community activities. I am also involved as a volunteer.
I have been a Headteacher in Surrey since September 1999, and am delighted to say that I am currently the Headteacher at North Downs. This is my third substantive Headship, and I have also acted as an interim Headteacher at two other schools, at the request of the local authority. North Downs is a very unusual school, in that we are spread over three sites in three different villages. However, the whole school community could not be stronger, with pupils, parents, staff and governors striving to "Be the best you can be, every day".
I have been associated with North Downs Primary school for a very long time. At first I was a parent helper, then went on to become a teaching assistant, and after gaining my degree became a teacher here through the Graduate Teacher Programme. I am currently a Year 3 teacher as well curriculum team leader.
This is my second term of office as a staff governor. In my spare time I have a variety of hobbies which include my love for reading as well as walking my dog.
I am a parent at the school as well as being privileged to be a co-opted governor.
Having previously helped at local toddler groups and the nursery, being a governor is now a great way to contribute my experience of managing staff and resources in small organisations to the school.
I currently work at Juniper Hall Field Studies Centre and look forward to working with the school on their outdoor learning.
I am delighted to be a Co-opted Governor of North Downs. I live in Epsom and have 2 children at primary school. I have had experience of being a Governor at a school in Ashtead and have taught for a number of years in Secondary Schools. However I now work in full-time ministry as an Associate Pastor of a church in Sutton.
Born and raised in Leigh I returned with my family in 2015. Our twin boys are happily attending Leigh school after a really enjoyable year at Betchworth.
With a background in the private sector mainly in finance as an accountant and systems as a director of a IT consultancy I'm keen to learn about and hopefully help with the opportunities and challenges that face the school.
I first came to Brockham more that 40 years ago when I started work as a Research Chemist at the Beecham Research labs at Brockham Park. I fell in love with the village and have lived in Strood Green for the last 30 years. I am now semi-retired and invest my time and energy in promoting chemistry at both public and education outreach events. As a Fellow and Trustee of the Royal Society of Chemistry I support their education and industrial strategies. I was honoured to be awarded the British Empire Medal in 2016 for voluntary service to the chemical sciences.
I still find time for my hobbies which include; natural history, gardening, walking baking, knitting and tatting.
I have lived in Brockham for nearly 10 years and have 2 children at North Downs Primary School. After having previously been involved in local committees in Brockham, I am delighted to be part of the Governing Body. I am totally committed to supporting the school in achieving the best possible outcome for all of its pupils.
I have an engineering background and currently work for Clarion Housing Group. In my spare time I love being outdoors and taking part in a variety of hobbies including running, wild swimming and walking.


text:
package com.raquo.dombuilder

import com.raquo.dombuilder.jsdom.simple.bundle
import com.raquo.dombuilder.jsdom.simple.bundle.{checked, div, input, option, selected, typ}
import com.raquo.dombuilder.jsdom.simple.implicits._

class PropSpec extends UnitSpec {

  it("sets props") {

    mount("input:checkbox", input(typ := "checkbox", checked := true))
    expectNode(input like (typ is "checkbox", checked is true))
    unmount()

    mount("input:text", input(bundle.value := "yolo"))
    expectNode(input like(bundle.value is "yolo"))
    unmount()

    mount("option=true", option(selected := true, bundle.value := "12"))
    expectNode(option like(selected is true, bundle.value is "12"))
    unmount()

    mount("option=false", option(selected := false, bundle.value := "12"))
    expectNode(option like(selected is false, bundle.value is "12"))
    unmount()

    mount("nested input:text", div("foo", input(bundle.value := "yolo")))
    expectNode(div like ("foo", input like(bundle.value is "yolo")))

    unmount()
  }
}

redpajama_set_name: RedPajamaGithub