As discussed in the previous sections, we finally obtained for each programme
star a good SED fit including accurate estimates
of the stellar astrophysical parameters, such as E(B-V), 
,
and 
values, which are collected in Table 7 (click here).
With the intrinsic stellar parameters 
and 
,
we can construct an HRD.
is simply estimated from the difference between the
observed V, for which we used the values from the best SED fits,
and 
, with 
.
As discussed in Sect. 5, for a few objects we have multiple
possibilities in the SED fit. They will be analyzed as being
separate objects. With the data of the remaining 52 objects we
construct an HRD in Fig. 4 (click here). In this figure we have discriminated
to 
values by plotting different symbols: a cross for
stars with 
= 3.1, a square for stars with 
3.1
and circles for objects with unknown extinction
characteristics: the objects which are denoted by sh.
Figure 4 (click here) also contains
two stellar birthlines of Palla & Stahler (1990 and 1992), which is the
boundary after which stars become optically visible when evolving towards
the main
sequence. The upper one is for an accretion rate on the protostar of
= 10
yr
and the
lower one for 
= 10
yr
. We have also drawn the Zero Age Main
Sequence of Hillenbrand et al. (1993).
The birthlines and ZAMS are recalculated from the published HRDs
to our observational HRD by adopting a distance of 2.6 kpc and
a luminosity class III. From Fig. 4 (click here) it is evident that the
distance modulus corresponding to a distance of 2.6 kpc seems
to be appropriate also for the PMS sample. We therefore
conclude that this estimate is correct.
Before discussing the cluster membership of the sample of objects, we will first analyze one of the main criteria for this: the HRD. We start with the assumption that all the programme stars are located at the same cluster distance.
Figure 4:
The 
, 
diagram for our programme stars
in NGC 6611. The data points plotted as 
are of stars with
= 3.1; those as 
with 
3.1 and those as
are for stars with an IR-excess probably due to circumstellar dust
radiation. The Groups I: 
0
1; II:
0
3
0
7; and III:
0
7 are clearly separated
From the location of our 52 PMS candidates in the HRD, Fig. 4 (click here), we can
separate three groups: Group I are all early-type stars,
0
1, and lie between the MS and one of the stellar
birthlines. Group II contain the stars of intermediate spectral types,
0
3 
0
7. They
are located near or just to the right of the upper
stellar birthline. Group III consists of the
late-type stars with 
0
7, which
are located far to the right and/or far above the stellar
birthline. We shall discuss now each group separately.
Group I
The early-type stars, the majority of our sample, fit nicely
to the MS as derived by #Th&Thé et al. (1990) or to the ZAMS of
Hillenbrand et al. (1993), or lie
somewhat to the right of it. However, located to the right of
the ZAMS we find the more luminous stars more frequently than
the lower luminosity ones. This can be easily understood if we
accept that several of these massive stars are already moving
away from the ZAMS, to become post-MS stars.
The early-type stars at the lower end of the MS are remarkable.
We find that 4 out of 5 of the early-type stars located in this
part of the HRD have a strong near-IR-excess, probably due to dust.
Although these objects are located near the ZAMS, they can still be
young enough to belong to the HAeBe group.
Furthermore, we see that for Group I stars there is no dependency of
the 
value for their location in the HRD.
Group II
From this very interesting group none seem to have ``normal'' IR properties,
either they have large 
values or they are probably surrounded by dust.
Although these objects are located to the right of the birthline for large
accretion rates one can not discard them as not being PMS if they are cluster
members. As shown by Palla & Stahler (1993) and in Fig. 4 (click here), it could be
possible that if we change the accretion rate the
location of a birthline could shift (it is discussed by Lamers et al.
1996, in preparation, that this can also happen by changing the metallicity of
a star).
For an accretion rate (
yr
) the Group II objects can be visible in their
PMS phase.
Group III
All stars in this group have a normal extinction law, are located far to
the right the birthline, are of very late spectral type and are relatively
luminous if we place them at the cluster distance. The objects in this group
are those objects for which Chini & Wargau (1990) concluded that late-type
stars are foreground objects. However, we have shown here (see Group II)
that we must discriminate between two types of late-type objects, depending
on their IR properties and their location in the HRD.
The exceptional but interesting location of the star at 
=
0
20, is the star W213(1). The classification of this star used in
option W213(2) is probably most appropriate one (see
Table 7 (click here) and Sect. 7.2).
From Fig. 4 (click here) we encounter a major problem: where are the stars that approach the ZAMS?, or in other words: ``where are the late A-type and F-type PMS objects?'' We return to this interesting problem later.
In order to make this discussion more conclusive, we first discuss the probabilities of cluster membership for the programme stars. After which we discuss the extinction characteristics of the cluster field. Subsequently, conclusions are drawn in combination with the properties of the HRDs.