rosat observations of cepheus ob3: the discovery of low-mass stars

ROSAT observations of Cepheus OB3: the discovery of low-mass stars

Tim Naylor1 and A. C. Fabian2

1Department of Physics, Keele University, Keele, Staffordshire ST5 5BG2Institute of Astronomy, Madingley Road, Cambridge CB3 OHA

Accepted 1998 September 25. Received 1998 September 21; in original form 1996 December 18

A B S T R A C T

We present X-ray images of the high-mass star-forming region Cepheus OB3, in which we

discover over 50 point sources, the majority of which are T Tauri stars. We use the ratio of high-

mass to low-mass stars to constrain the initial mass function, and ®nd that it is consistent with

that for ®eld stars. This supports the picture that there is little difference between high-mass

and low-mass star-forming regions. The presence of low-mass stars between the high-mass

stars and the molecular cloud they are ionizing argues against mechanisms whereby heating of

the cloud inhibits low-mass star formation. Finally, the ef®ciency of star formation in this

region is shown to be of the order of a few per cent.

Key words: stars: formation ± stars: luminosity function, mass function ± stars: pre-main-

sequence ± open clusters and associations: general ± open clusters and associations:

individual: Cepheus OB3 ± X-rays: stars.

1 I N T R O D U C T I O N

The observational distinction between the high- and low-mass star-

forming regions is gradually being eroded. The Orion OB associa-

tion, traditionally viewed as a high-mass star-forming region, has

long been known to harbour low-mass T Tauri stars, and Walter et al.

(1994) have discussed the low-mass stars in the Upper Sco associa-

tion (Sco OB2-2). Recently, Schulz, BerghoÈfer & Zinnecker (1997)

and Gregorio-Hetem et al. (1998) have also presented X-ray

observations of distant OB regions which seem to show the

presence of T Tauri stars. The data presented here con®rm this

view, with the discovery of low-mass stars in another OB region,

Cepheus OB3. Walter & Boyd (1991) have shown that the com-

plementary case, that low-mass, or T associations, should contain

high-mass stars is probably correct, at least for Taurus-Auriga.

If there is no physical distinction between these two types of star-

forming region, then it would seem natural to assume that there is

just one star formation mechanism which operates within them. The

strongest evidence in favour of such a mechanism would be a

universal initial mass function (IMF), although if there were

different mass functions, this would not necessarily rule out a

single star formation mechanism. Despite the large number of mea-

surements now available, there is no strong evidence against a

ubiquitous IMF for stars greater than 1 M(, although the evidence

becomes weaker as the IMF becomes more structured below this mass.

Furthermore, measurements that span a large range of mass, particu-

larly those that cover both the late-type and early-type stars, are rare.

2 T H E C E P H E U S O B 3 A S S O C I AT I O N

The Cepheus molecular cloud is clearly visible on sky survey plates

as a region of low stellar density. The outlines traced by the stars

match the J � 1 ! 0 12CO maps of Sargent (1977, 1979) well. The

molecular cloud is structured, with a strong CO peak at its most

north-western point, known as Cepheus B. Outside Cepheus B, to

the north and west of the cloud, lies the stellar association Cepheus

OB3 (see Fig. 1). It is divided by Blaauw (1964) into two popula-

tions, OB3b being younger than OB3a. These ages are, however, the

subject of some controversy, since the turn-off ages are 7.5 and

5.5 Myr (Jordi, Trullols & GaladõÂ-EnrõÂquez 1996), whilst the

expansion ages are 0.72 and 0.48 Myr (Garmany 1973).

There are obvious signs of star formation activity on the edge of the

molecular cloud closest to the OB3b association. This north-western

edge of the cloud is bounded by an optical H II region Sh 2-155

(Sharpless 1959), in which material is being driven off the cloud by

the O7n star HD 217086 (Panagia & Thum 1981). A trapezium of

stars appears to have recently emerged from the edge of the cloud

(Testi et al. 1995), and Cepheus B itself is thought to be in the early

stages of contraction towards becoming a new cluster of stars.

Perhaps the most important aspect of Cep OB3 is the very

favourable viewing angle at which we see the progression of star

formation. As the surface of the cloud is being eroded by the

radiation from one or more of the early-type stars, the cloud edge

moves eastwards across our ®eld of view, with the newly emergent

stars clear of the molecular cloud and to the west. This is in contrast

to the rather awkward view we have of the Orion OB1 region, where

the stars lie between the cloud and ourselves.

Throughout this paper we will use the distance and reddenings

for Cep OB3b derived by Moreno-Corral et al. (1993) of 900 6 100

pc (which is consistent with most measurements) and AV � 2:9.

3 O B S E RVAT I O N S

We observed a ®eld centred on the O7n star HD217086 with both

Mon. Not. R. Astron. Soc. 302, 714±722 (1999)

q 1999 RAS

the ROSAT PSPC and HRI, a log of the observations being given in

Table 1. Fig. 2 shows the two PSPC images, and the HRI image. The

bright source at the centre of all the observations is the O7n star,

which lies in the northern section of what appears to be a cluster of

X-ray sources. In addition to the point sources, there is underlying

unresolved diffuse X-ray emission. The Einstein IPC observation of

this ®eld (Fabian & Stewart 1983) detected the X-ray emission from

this cluster, but was unable to resolve it into individual sources.

There are virtually no X-ray sources to the south-east, which is

unsurprising as this is the region of the molecular cloud (see Fig. 1).

There are, however, a few sources to the north and east of the main

cluster. The variability of several sources between the two PSPC

observations is obvious, the most spectacular example being the

source to the north-east of the O7n star (Star 45 in Table 2), which,

aside from the O7n star, is the brightest object in the ®rst PSPC

observation, and yet is barely detected in the second (see also

Table 2). There is almost certainly also variability between the

PSPC and HRI observations, but the different passbands of the two

instruments should make us cautious about such an interpretation of

the data on a source-by-source basis.

We ®rst created a catalogue of all the X-ray sources we detected

by searching the central 40-arcmin diameter of the HRI and two

PSPC images using the point-source search algorithm of Allan

(1992), which is also described brie¯y in, for example, Jeffries,

Thurston & Pye (1997). We searched to a level of 4.5j at a given

position in the PSPC images, and 4.25j in the HRI images, which

simulation shows should result in less than one spurious source per

image (Jeffries et al.. 1997; Jeffries, private communication). Each

of the source lists contained at least three objects that could be

identi®ed with stars in the HST GSC. This allowed us to carry out a

bore sight correction, and sum the two PSPC images, so that a

further search for sources could be carried out. As a result of this,

the maximum error in either RA or declination between an X-ray

position and a GSC position was 15 arcsec for the combined PSPC

image, and 7 arcsec for the HRI. The bore sight correction was

applied to each list, before attempting to combine them to produce

a single catalogue for the whole association. Since many of the

sources are variable, we searched for positional coincidences

amongst the lists, ignoring the ¯ux levels of the objects. Sources

within 15 arcsec were considered to be the same object. This

separation is suf®ciently large that it may occasionally result in

two sources being assumed to be one object, when in fact they are

unrelated. However, in the statistical analysis presented later, it is

better that this occurred, than that a single object should be

represented twice in our catalogue.

Once the lists had been merged, the coordinates for each object

were taken from the best available source, i.e., the HRI where

available, otherwise the combined PSPC image or, in two cases, the

single PSPC images. These positions are presented in Table 2. We

then evaluated the ¯ux at each of these positions in all three

observations, correcting for the gross vignetting, and in the case

of the PSPC, the vignetting due to both the thin and thick wires.

Where the ¯ux was either unmeasurable with the point-source

search algorithm, or below the error, we calculated a 1j upper

limit. Both the ¯uxes (with 1j errors) and limits are presented in

Table 2.

With these ¯uxes, we carried out a more formal variability

analysis. First, we looked for sources that varied between the two

PSPC observations at greater than the 99 per cent con®dence level,

®nding that Sources 21, 22, 23, 36, 44 and 45 did so. Next, we

searched for variability within each observation for the observations

of sources which had more than 50 detected photons, excepting

Sources 6 and 7 which are too close to allow reliable photometry.

For each source we extracted a light curve in bins of 1000 s for the

PSPC and 2000 s for the HRI, and then tested to see if the resulting

light curves were consistent with a constant ¯ux using a x2 test.

Only Source 23 in PSPC1 was inconsistent at a greater than 99 per

cent con®dence level. This, and the null results of our other tests, are

reported in Table 2. Since the total number of photons detected in

the PSPC1 observation of Source 23 is small, all one can say with

certainty is that the source is brighter at the beginning of the

observation than at the end.

To be certain of the nature of these sources, we required spectro-

scopy of their optical counterparts. To generate candidates for

spectroscopic observations, we ®rst checked the positions against

the HST GSC. We report the 10 positive results in Table 2. The

modal magnitude (in 1-mag bins) of the HST GSC in this region is

13. Thus we expect that the catalogue of stars outside the region of

our CCD data is complete to about 12th magnitude.

The optical ¯ux limit from the GSC is above the ¯ux we would

expect for T Tauri stars, and so we required deeper imaging of the

®elds. Thus we observed the region using the 0.6 m Thornton

re¯ector at Keele Observatory during the night of 1995 December

29. The system uses an ST6 CCD camera at the f =4:5 Newtonian

ROSAT observations of Cepheus OB3 715

q 1999 RAS, MNRAS 302, 714±722

Figure 1. A schematic of the Cepheus OB3b region. The dashed line marks

the lowest CO contour, taken from Sargent (1977); the molecular cloud is to

the lower left of this line. The solid line is the brightest part of the Ha nebula

Sh 2-155, and is another way of marking the edge of the cloud. The open

circles are the early-type members of the Cepheus OB3b association taken

from BHJ and Sargent (1977), with the O7n star BHJ 41 marked. The crosses

are the sample of coronal X-ray sources discussed in the text. The dotted

circle is the PSPC inner support ring, and represents the extent of the area

searched for X-ray sources.

Table 1. Log of ROSAT observations.

Instrument Date start Date end Total exposure Bore sight correction

PSPC 1992 July 29 1992 July 30 4252 s �1:s0, ÿ0500

PSPC 1992 December 27 1993 January 6 6357 s ÿ1:s1, �0300

HRI 1994 July 21 1994 July 24 20308 s ÿ0:s7, ÿ0100

focus, and in our case we used a Johnson V-band ®lter. We observed

four ®elds for two min each, covering a region of approximately 14

by 17 arcmin2, centred on RA 22h 56:m7, Dec. �628410 (J2000). The

images were dark-subtracted, and combined into a single image

that we could use for astrometry. We searched for positional

coincidences between our X-ray detections and the optical ®eld.

We chose to accept correlations within 7 arcsec of an HRI position,

and 15 arcsec of a PSPC position, irrespective of whether the

counterparts were from the GSC or our own data. The results are

listed in Table 2.

We measured optical magnitudes for the counterparts using the

aperture photometry algorithm described in Naylor (1998). In

addition to the long exposures, we obtained a single 10-s exposure

that had stars in common with all four long exposures, allowing us

to transfer all the data to a single zero-point. Since the short

exposure also contained both Stars 40 and 41 of Blaauw, Hiltner

& Johnson (1959, hereafter BHJ), we used their photometry of these

stars to calculate our overall zero-point. We have not attempted any

colour correction, but the correction for this system is comparable

with our quoted accuracy (0.1 mag). The faintest objects visible in

the optical data correspond to about V � 16:8, which is an approxi-

mately 10j detection. We will use this as the upper limit to the

optical ¯ux for those sources which have empty error boxes. The

chance of a random correlation between a star brighter than

V � 16:8 in our CCD images and, say, a 10-arcsec radius error

circle is highly dependent on position within the association,

because the stellar density changes dramatically across the ®eld-

of-view. For the densest parts of the cluster, however, it is around 10

per cent. Thus we expect one or two of our 16 identi®cations based

on the CCD imaging alone are spurious.

716 T. Naylor and A. C. Fabian

q 1999 RAS, MNRAS 302, 714±722

Figure 2. The X-ray images from observations PSPC1 (upper left), PSPC2 (upper right), HRI (bottom left) and PSPC2 overlayed on the Digitised Palomar Sky

Survey (bottom right). The lowest contour is at 0.1 counts per 15-arcsec pixel. Each subsequent contour is at a factor of 2 greater counts than the last. The PSPC

images have been smoothed with the adaptive routine described in Ebeling, White & Rangarajan (1997), and are for the 0.5±2 keV band.

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Optical spectra of 10 of the tentative identi®cations were obtained

for us at the 2.5-m Isaac Newton Telescope (INT), during the night

of 1995 October 7/8. The Intermediate Dispersion Spectrograph

was used with the 235-mm camera, R1200R grating and Tek CCD.

This covered the range 6140±7070 AÊ at a resolution of 1.7 AÊ and a

dispersion of 0.85 AÊ pixelÿ1. The spectra were debiased using an

overread region of the chip, and ¯at-®elded using an exposure of a

tungsten lamp. The spectra were extracted using the procedure

described in Horne (1986), and wavelength-calibrated using exposures

of arc lamps.

4 W H AT A R E T H E P O I N T S O U R C E S A N D

D I F F U S E E M I S S I O N ?

The obvious interpretation for the majority of the point sources is

that they are T Tauri stars in the association. The discovery of T

Tauri stars via their X-ray emission is now commonplace (see

NeuhaÈuser 1997 for a review), and, as we show below, the optical

and X-ray luminosities of these sources match this hypothesis. First,

however, we consider the likely levels of contamination from

foreground and background stars, and from background AGN.

Guillout et al. (1996) present a model for the stellar content of

ROSAT observations, which Motch et al. (1997) evaluate for

galactic latitude b � 08 and longitude l � 908, yielding about two

sources brighter than 5 counts per kilosecond. This can be com-

pared with the approximately 40 sources found in our combined

PSPC image above this count level. Furthermore, our pointing is at

about b � 38 and l � 1108, both of which changes in coordinate

tend to reduce the number of contaminating sources. A hard upper

limit to the number of AGN in our data can be derived using the

ROSAT Lockman Hole observations presented in Hasinger et al.

(1993, 1994). From their log(N)±log(S) function there would be

about three sources with ¯uxes above 4 count ksÿ1 in the 0.5±

2.4 keV band, again compared with our approximately 40 sources

found above this ¯ux in the combined PSPC image. In fact, Galactic

absorption should reduce this number considerably, but precise

calculation of this reduction is hard, due to the highly structured

nature of the plane in this region. However, the conclusion is clear:

there is a signi®cant excess of sources in this region above that

expected for foreground and background objects.

To examine the X-ray and optical luminosities of the sources, we

®rst constructed the diagram of optical versus X-ray ¯ux, shown in

Fig. 3. Studies of ®eld stars show that most single stellar X-ray

sources are either earlier than spectral type A5, with X-ray emission

associated with their stellar winds, or later than F5, with coronal

X-ray emission. Hence we expect the upper part of Fig. 3 to be

occupied by X-ray-bright O and B stars, and the lower part to be

®lled with coronally active stars. We have therefore coded the

points plotted in Fig. 3 according to spectral type. Previously

discovered O, B and early A stars, and one discovered via our

INT spectroscopy, are marked as circles. Since our INT spectra

would detect the strong Ha absorption characteristic of B and early

A stars, we mark the remaining stars for which we have spectro-

scopy as late-type, with open triangles; if they show Ha emission,

the triangles are ®lled. Stars for which no spectral-type information

is available are marked with crosses. Stars 47 and 48 have been

excluded (see Section 8). Fig. 3 shows that we do see a clear

boundary between the early-type stars (circles) and the coronal

emitters (triangles), at V � 12:5.

We can determine the X-ray luminosity only if we assume a

spectral shape, and so we ®tted the spectrum of Star 34 to a

Raymond±Smith model. The best-®tting temperature is .1 keV.

Using temperatures in the range 1 to 3 keV, typical of T Tauri stars,

gives a conversion of 10 count ksÿ1 which corresponds to 1031:1

erg sÿ1 (1±2.4 keV). Comparing this with the T Tauri stars in

Chamaeleon I (Feigelson et al. 1993) shows that these objects lie at

the top of the T Tauri X-ray luminosities. Although there is

controversy about T Tauri luminosity functions (e.g. NeuhaÈuser


q 1999 RAS, MNRAS 302, 714±722

Figure 3. The V-band magnitude and X-ray ¯ux of all the sources detected

in the ROSAT observations. The different symbols represent different

spectral types. Stars later than mid-A are shown as either open triangles,

or, if they have emission, ®lled triangles. O, B and early A stars are open

circles, and those with no spectral information are crosses.

Figure 4. The INT spectra of the late-type stars. Note the lithium absorption

at 6708 AÊ in some of the spectra, and the Fe I 6595-AÊ blend. Some of the

spectra have absorption or emission features at Ha (6553AÊ ), and all show

telluric absorption at about 6870AÊ .

et al. 1995, and references therein), especially the difference

between weak and classical T Tauri stars, this does not affect

what the highest luminosities are that one might expect to see.

The T Tauri interpretation is given further weight by the ratio of

X-ray to V-band luminosity for the optically faintest sources, which

is ,10ÿ1:5 as one might expect for very active stars (e.g. Stocke et al.

1983).

Several of our spectra show lithium 6708 AÊ at impressive

(<300 mAÊ ) equivalent widths (see Table 2 and Fig. 4). Whilst

this is normally thought of as an indicator of mid- or late-type pre-

main-sequence stars, weaker lithium also occurs in young main-

sequence stars. Since the lithium equivalent width is a function of

spectral type for a given abundance, spectral types are needed to

proceed further. We therefore adopt a simpli®ed version of the

approach of Magazzu et al. (1997). We examined the spectra of the

six stars which show unblended Li absorption, and compared them

with the spectral atlas of Jacoby et al. (1984). None of them shows

TiO features, and so they are earlier than late K. By examining other

features, especially the equivalent width of the Fe I 6595-AÊ blend,

we identi®ed two stars as G, and four as K. Soderblom et al. (1993)

show that the Pleiades have equivalent widths of about 0.2 AÊ for G

stars, 0.4 AÊ for early- to mid-K stars, after which the equivalent

width falls to almost zero. Stars 32 and 34 are K stars with Li

equivalent widths of 0.50 and 0.38 AÊ , and are therefore younger

than the Pleiades. The remaining stars are either close to the

Pleiades equivalent widths (G stars 23 and 33), or depend critically

on spectral type (K stars 13 and 46), which is hard to determine from

our low signal-to-noise spectra. However, it is clear that some of the

stars are pre-main-sequence, based on their Li absorption.

Two of the eight sources for which we have spectra clearly show

Ha emission, although only one of them is above the formal 10-AÊ

limit normally taken as the de®nition of a classic T Tauri star. This

would seem to indicate that there are more weak than classical T

Tauri stars in the luminosity range explored by our observations.

However, such a conclusion is at best preliminary, since we does not

have optical spectroscopy for a complete sample above a given ¯ux

limit.

Summarizing the above arguments, the point sources are

certainly not active ®eld stars or AGN. Their X-ray and optical

luminosities are consistent with T Tauri stars, as is the ratio of X-ray

to V-band luminosity. Many of the optical spectra show strong Li

6708-AÊ absorption, in some cases strong enough to be almost

certainly pre-main-sequence stars, and some show Ha emission.

Finally, the X-ray variability is similar to that observed in other

groups of T Tauri stars (e.g. Montmerle et al. 1983). For these

reasons, it seems likely that we have discovered low-mass stars in

Cep OB3.

To examine the nature of the diffuse X-ray emission within the

cluster, we compared its spectrum with that of the point sources. We

®rst excised the point sources from within the cluster, and then

created a spectrum from the remaining counts. We used a region to

the south-west, i.e., in the direction of the molecular cloud, as

background. We then compared this with the background-

subtracted sum of the point sources. The resulting spectra, shown

in Fig. 5 are indistinguishable. Given that the diffuse emission also

lies within the region of highest density of point sources, it is likely

that it is simply unresolved T Tauri stars.

5 T H E S PAT I A L D I S T R I B U T I O N O F T TAU R I

S TA R S

If we wish to study the spatial distribution of the T Tauri stars, we

must ®rst create a clean sample. For simplicity, we dealt only with

those X-ray sources that were detected in the combined PSPC

image. The very faintest sources also present a problem, because the

PSPC vignetting gives a lower ¯ux limit for sources in the centre of

the ®eld than at the periphery. (The background is low enough in

10 ks that the increase in point spread function does not affect the

detection threshold.) The vignetting is almost 20 per cent at 20

arcmin off axis, and so we chose to exclude those sources whose

¯ux was within 20 per cent of the faintest detected source, i.e., less

than 5.4 count ksÿ1. Finally, we wished to reject all the early-type

stars, and so applied the magnitude criterion discussed in Section 4,

by rejecting all those X-ray sources with optical counterparts

brighter than V � 12:5. Since the HST GSC is complete to approxi-

mately this magnitude, we can be relatively sure that we have

detected optical counterparts to this magnitude even outside the

®eld covered by our CCD images. If a high-mass star happened to

lie within the error box of a low-mass X-ray source, our technique

would reject it from the sample. However, with just eight high-mass

stars within the inner circle of the PSPC, and using a 7-arcsec X-ray

error circle, there is only a 1 per cent chance that a low-mass X-ray

source has been accidently rejected in this way. We also rejected

Source 47, for reasons outlined in Section 8.

In Fig. 1 we show the resulting source positions, along with the

edge of the molecular cloud, the edge of the PSPC central ®eld, and

the positions of the OB stars which are members of the Cep OB3b

association. The latter are those originally identi®ed by BHJ, with

the membership criteria re®ned by Sargent (1977). Assuming the

Ha emission de®nes the edge of the cloud, there are very few

sources within the cloud itself. Outside the cloud the sources do not

appear to be uniform in density, but cluster close to the cloud edge.

Assessing the signi®cance of the density enhancement is very

dif®cult, since it lies on a rapidly varying background of objects.

However, it is associated with the diffuse background discussed

above, and so it seems likely that the density of T Tauri stars is

highest close to (but outside) the cloud edge. This is in contrast to

the OB stars in the region, which seem to be strung out slightly

further from the cloud.

Before attaching much signi®cance to the impression that the

T Tauri stars are closer to the cloud than the high-mass stars, we

performed a two dimensional Kolmogorov±Smirnov test (Press

et al. 1992) to see if the distribution of T Tauri stars and OB stars


q 1999 RAS, MNRAS 302, 714±722

Figure 5. The PSPC spectra of the diffuse background (dots) and the point

sources (stars). Within the error bars the two spectra are identical, consistent

with the idea that the diffuse X-ray emission is just unresolved point sources.

was actually different. For this test we require a well-behaved

sample of massive stars, and so we used those observed by BHJ and

classi®ed as members by Sargent (1977). These simple criteria are

possible, since BHJ observed a sample complete to approximately

12th magnitude. The Kolmogorov±Smirnov test gave an 41 per

cent probability that both samples were drawn from the same parent

spatial distribution. So we conclude that although it may appear that

there is a cluster of T Tauri stars on the outside edge of the cloud, we

cannot prove that they have a signi®cantly different distribution

with respect to the OB stars.

Our X-ray observations of Cep OB3 are similar to those of

IC 1396 (Schulz et al. 1997) and the Monocerous R2 and Rosette

nebulae (Gregorio-Hetem et al. 1998). All these observations show

many point sources, which are almost certainly at the tip of the T

Tauri luminosity function, and diffuse emission which can be

interpreted as unresolved T Tauri stars. There is one interesting

difference, however. From Fig. 1 it is clear that, with a few

exceptions the Cep OB3 sources avoid the area of CO emission.

In contrast, the sources in Mon R2 and the Rosette seem to be mixed

with the CO (see especially ®g. 13 of Gregorio-Hetem et al.). Again,

in Orion Sterzik et al. (1995) ®nd T Tauri stars embedded within the

cloud, in addition to being distributed outside it. It seems that, in

contrast, Cep OB3 is particularly simple in its star formation

history. The ionization front appears to be moving from west to

east, resulting in a group of newly formed stars to the west of the

cloud edge. Presumably a deeper exposure would reveal more

sources in the molecular cloud itself.

The simplicity of Cep OB3 does allow us to make one very

important inference, however. It is clear that low-mass stars have

emerged from the molecular cloud despite it being strongly irradi-

ated by nearby early-type stars. This allows us to rule out any

suggestion that high-mass stars beget only more high-mass stars

(e.g. Larson 1986), inhibiting low-mass star formation by raising

the temperature of the molecular cloud.

6 T H E I N I T I A L M A S S F U N C T I O N

We can constrain the IMF between about 1 and 17 M( using the

ratio of the number of detected X-ray sources to the number of OB

stars within the PSPC ®eld. Our total number of detected low-mass

X-ray sources is about 50, which represents an absolute minimum to

the number of stars in the T Tauri mass range of 0.2 to 2.3 M(

(Cohen & Kuhi 1979); however, this sample is incomplete because

of the (variable) ¯ux limit. The number of high-mass stars in the

region covered by the PSPC is 8, of which the lowest mass is

approximately 4 M(, based on the apparent magnitude limit of

V � 12. Thus the ratio 0:3 < M < 3 M( : M > 4 M( is certainly

more extreme than 6:1. Further, in the region of the cluster of

sources, there is roughtly equal ¯ux from resolved and unresolved

sources. To be unresolved these sources must, individually, be

fainter than the resolved sources. Thus a lower limit to the

number of sources represented by the unresolved ¯ux is that they

are equal in number to the resolved sources. So we can tighten our

limit on the ratio of low-mass to high-mass stars to be more than

12:1. Perhaps the most realistic ®eld IMF to compare this with is

that of Miller & Scalo (1979), which gives a ratio for the same mass

ranges of 25. Two other possible mass functions are those due to

Salpeter (1955), which would require a ratio of 30:1, and Van Buren

(1985), which gives 13:1. So we conclude that although normally

regarded as a massive-star formation region, Cep OB3 certainly

contains low-mass stars, and our lower limit on the number of them

is not inconsistent with ®eld IMFs.

The only other regions for which similar experiments have been

carried out are Orion, the Upper Sco Region (Sco OB2-2) and,

perhaps more importantly, the `low-mass' Taurus-Auriga region.

Unfortunately the Orion IMF has been examined only in the range

3 to 9 M( (Claudius & Grosbùl 1980), and so is not a useful

comparison. Walter et al. (1994) showed that the IMF of the Uppper

Sco Region is, like that of Cep OB3, consistent with a Miller &

Scalo IMF for masses greater than 0.3M(. The data for the Taurus-

Auriga IMF at high masses are more speculative (Walter & Boyd

1991), but are consistent with the Salpeter, Miller & Scalo and Van

Buren IMFs. Thus the data collected to date suggest that the IMFs

over quite a wide range of masses are similar in low-mass and high-

mass star-forming regions. Furthermore, it is clear that there is no

evidence for a de®cit of low-mass stars in OB regions when

compared with the ®eld.

7 T H E S TA R F O R M AT I O N E F F I C I E N C Y

We can estimate the star formation ef®ciency, i.e., what fraction of

the mass of the molecular cloud becomes stars, on the assumption

that the molecular cloud will continue to form stars at roughly the

same density as it has in the recent past. To do this, we compare the

mass of stars per unit area outside the cloud with the surface density

(on the sky) of the cloud itself. Minchin, Ward-Thomson & White

(1992) estimate that there are ,6 M( per square arcminute in the

`Cepheus B' region, which is the bright section of the molecular

cloud roughly enclosed by the nebulosity marked in Fig. 1.

The total mass of stars within the PSPC ®eld-of-view can be

found by integrating the mass for the Miller & Scalo IMF. This

yields an average mass (i.e.,�

MdN=�

dN) of 0.64 M(. The total

number of stars can be obtained by extrapolating the same IMF

from the eight high-mass (>4 M() stars, which yields about 200

stars in total. The area examined is certainly no larger than the 40-

arcmin diameter of the central region of the PSPC (1260 arcmin2),

but we will reduce this by 50 per cent, as the area in the eastern

half of the ®eld-of-view is either part of the molecular cloud, or

has not been involved in the recent star formation. This yields a

mass estimate of 0.2 M( per square arcminute. Assuming that this

region will form a similar density of stars to that close to the

molecular cloud, this implies a star formation ef®ciency of 3 per

cent.

This is a factor of 10 higher than that found in the l Ori region by

Duerr, Imhoff & Lada (1982), although star formation in that OB

region may have been truncated by irradiation from an OB wind. It

is closer to the 1±2 per cent ef®ciency found for the low-mass star-

forming clouds in Taurus by Mizuno et al. (1995). However, all

these estimates are probably suf®ciently uncertain that the only

conclusion that can be drawn is that Cep OB3b is not unusual when

compared with other star-forming regions.

8 T H E ` T R A P E Z I U M '

Sources 47 and 48 coincide with a trapezium of stars located on the

edge of the molecular cloud, in the midst of the Ha nebulosity. The

HRI position (47) coincides with the easternmost of the stars,

although it is possible that some of the emission could originate

from elsewhere. However, we can de®nitely rule out the suggestion

by Testi et al. (1995), made on the basis of the current PSPC data

alone, that the X-ray source is coincident with their compact radio

source `B'.

The presence of a clearing nebulosity suggests that the stars

suffer a large extinction and/or are pre-main-sequence even at early


q 1999 RAS, MNRAS 302, 714±722

spectral types. Such a conclusion is supported by the colour±

magnitude diagram for this region of Moreno-Corral et al. (1993).

Even though we have long-slit spectra, the diffuse Ha in this region

is changing so rapidly it is dif®cult to subtract a reliable back-

ground. However, the spectrum of the northernmost star does

appear to have strong Ha absorption, supporting the idea that

these are heavily reddened early-type stars emerging from the

molecular cloud; they are perhaps the beginnings of the next

generation of stars in Cep OB3. Given its possible early-type

nature, we excluded Sources 47 and 48 from our sample of late-

type stars.

9 C O N C L U S I O N S

There seems little doubt that the ROSAT pointings have revealed the

coronally active low-mass stars in Cepheus OB3b. Studying these

has allowed us to draw three main conclusions.

(1) The presence of low-mass stars between high-mass objects

and the molecular cloud is dif®cult to reconcile with the idea that

ionization from early-type stars inhibits low-mass star formation

(Larson 1986).

(2) The ratio of high to low-mass stars is consistent with the

initial mass function of ®eld stars. This means that Cep OB3b ®ts in

with the idea that there is no difference in the IMF between high-

mass and low-mass star-forming regions, furthering the case that

there is no observational distinction between the two classes of star-

forming region.

(3) The star formation ef®ciency in Cep OB3b is of order a few

per cent.

AC K N OW L E D G M E N T S

TN acknowledges the support of a PPARC Advanced Fellowship,

and ACF thanks the Royal Society for support. Some of the data

reduction was carried out on the Keele Starlink node using the ARK

software. We are grateful to those who support Keele Observatory,

especially the local amateur astronomers who help maintain the

telescopes, and ICL who have provided computing facilities. We

thank Mark Somers, Coel Hellier and Gregory Beekman who

obtained the spectra for us, and Rob Jeffries who commented on

the manuscript. The Isaac Newton Telescope is operated on the

island of La Palma by the Royal Greenwich Observatory in the

Spanish Observatorio del Roque de los Muchachos of the Institutio

de Astrophysica de Canarias.

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This paper has been typeset from a TEX=LATEX ®le prepared by the author.


q 1999 RAS, MNRAS 302, 714±722

rosat observations of cepheus ob3: the discovery of low-mass stars

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