Already Burbidge & Burbidge (1956, ApJ, 124, 116) made the first quantitative abundance analysis of several λ Bootis stars and found metal deficiencies by a factor of 20 relative to the Sun.

The next two decades brought only little additional information. Sargent (1965, ApJ, 142, 787) suggested in λ Bootis stars that nuclear reactions have destroyed most of the heavier elements in the outer layers of the atmosphere in λ Bootis stars. The stellar interior should therefore be helium rich and highly evolved (already beyond the red-giant phase). Kodaira (1967, PASJ, 19, 556) was the first who noticed that the oxygen abundance for a sample of λ Bootis stars was normal, something which was later verified by Baschek & Searle (1969, ApJ, 155, 537). This immediately ruled out the nuclear reaction hypothesis proposed by Sargent (1965).

Slettebak et al. (1968, AJ, 73, 152) used, for the first time (later also used by Hauck & Slettebak, 1983, A&A, 127, 231), the space velocity as well as the "moderately large rotational velocity" to distinguish λ Bootis from true Population II type stars.

The early eighties saw a new flood of papers dealing with λ Bootis stars. With the work of Cowley et al. (1982, ApJ, 254, 191) the topic of the λ Bootis phenomenon was raised again. They found abundance variations of ±0.3dex for iron in a sample of 34 "normal" late B and early A-type stars. Considering their sample to be representative, 10 to 20% of all "normal" type stars should show underabundances. Thus, the metal deficiency of the λ Bootis stars would just be a natural variation within the variety of all "normal" type stars. Various groups now tried to establish criteria, which allow the compilation of a larger homogeneous list of λ Bootis stars. Hauck & Slettebak (1983) investigated the group properties in the Geneva and Strömgren photometric systems. Abt (1985, ApJS, 59, 95) used classification resolution spectra to present an extensive list of probable candidate stars. Based on this knowledge, Hauck (1986, A&A 154, 349) established criteria to select new λ Bootis type candidate stars via photometric indices. At the same time, Baschek et al. (1984, A&A, 131, 378) used IUE spectra to investigate the group properties in the UV region. The λ Bootis stars were slowly recognized as a separate group in the family of A-type stars as shown by the review on chemically peculiar A-type stars by Faraggiana (1987, Ap&SS, 134, 381).

A mile stone was set by Gray (1988, AJ, 95, 220). He tried to homogenize the group definition based on spectroscopic as well as on photometric information. He formulated a working definition by which he was able to establish a group of 16 stars with very similar properties. Furthermore this definition allowed to distinguish the λ Bootis stars from other type of stars (e.g. true Population II, Field-Horizontal-Branch, He-weak, etc.).

It is important to note in the context of now upcoming theories that Bohlender & Landstreet (1990, MNRAS, 247, 606) found no significant magnetic field for a sample of λ Bootis stars.

The λ Bootis stars are early-A to early-F type stars, with an approximate spectral type range (based on the hydrogen lines) of B9.5 to F0 with possible members as late as F3.

The λ Bootis stars seem to be always characterized by weak MgII 4481 lines, such that the ratio MgII 4481/FeI 4383 is significantly smaller than in normal stars. In addition the spectra exhibit a general metal-weakness. The typical shell lines (FeII 4233, etc.) tend to be weak as well but the λ Bootis stars do not show the typical shell spectral characteristics.

The following classes of stars should be excluded from the λ Bootis even if they show weak λ4481 lines: shell stars, protoshell stars, He-weak stars (easily distinguished on the basis of their hydrogen line temperature types), and other CP stars. FHB and intermediate Population II stars may be distinguished from the λ Bootis stars on the basis of their hydrogen line profiles. High-v.sini stars should be considered as λ Bootis candidates only if the weakening of λ4481 is obvious with respect to standards with high values of v.sini.

λ Bootis stars are also characterized by broad hydrogen lines, and in many cases, these hydrogen lines are exceptionally broad. In the late-A and early-F λ Bootis stars, the hydrogen line profiles are often peculiar, and are characterized by broad wings, but shallow cores.

The distribution of rotational velocities of the λ Bootis stars cannot be distinguished from that of normal Population I, A-type stars.

This definition already includes the basic criteria for a successful classification of a true λ Bootis star. The above given criteria limit the useful dispersions for classification surveys from 40Å/mm to 120Å/mm. Dispersions lower than 120Å/mm make it difficult to detect the characteristics of the λ Bootis phenomenon. The broad hydrogen line profiles are difficult to normalize properly without a good expanse of the continuum on either side of the line, and this leads to a limit for the highest dispersion of 40Å/mm.

This figure shows the classification resolution spectra of four well established λ Bootis type objects together with appropriate MK standards. The weakness of the metal-lines is always coupled with the weakening of the CaII K line. The notation of the spectral types is according to Gray (1988).

The classification process is not always straightforward. Different classifiers have established or rejected a group membership for several individual objects. Faraggiana & Bonifacio (1999, A&A, 349, 521) and Paunzen (2000, PhD thesis, University Vienna) list an extensive summary of these objects. Another highly underestimated effect is the influence of fast rotation on the estimated spectral type. As Slettebak et al. (1980, ApJ, 242, 171) have shown that the equivalent width of this line decreases with increasing rotation for models later than A0. This fact was also described by Abt & Morrell (1995, ApJS, 99, 135). A much stronger effect was found for Hγ. This means that a fast rotating star will be classified much later using Hγ than using MgII 4481. Taking a rapidly rotating A5 type star one would classify it as: hF1mA7 or metal-weak.

We (1997, A&AS, 126, 407; 2001, A&A, 373, 625; 2001, A&A, 373, 633) have presented an extensive classification resolution spectroscopic survey in the Galactic field and nine open clusters in order to detect new λ Bootis stars. They classified 708 stars which were selected on the basis of narrow band photometry. They confirmed the membership of 18 candidates and detected 26 new λ Bootis stars. Furthermore, 44 metal-weak objects were found which might be connected to the λ Bootis group. Another long-term survey was devoted to members of open clusters by Gray & Corbally (2002, AJ, 124, 989). They noticed the apparent absence of λ Bootis stars in 24 open clusters with ages from 15 Myr to 700 Myr.

The c₁ and m₁ versus (b - y)₀ plots for λ Bootis type (filled circles) and "normal" (open circles) stars are equivalent to the corresponding diagrams in the Geneva 7-color system. From this figure we are able to conclude:

Metallicity: All members have a low metallicity which decreases with temperature. A photometric parameter space which includes all candidate stars is well determined.

Temperature: The temperature ranges from late B V to early F V stars.

Surface gravity: λ Bootis stars cannot be distinguished from normal dwarf stars

Confusion with non−λ Bootis stars: It is obvious that also other stars populate the same parameter space. An unambiguous detection of λ Bootis stars with standard photometry alone is not always successful, although to some extend a discrimination with metallicity sensitive indices is useful (e.g. the Δa index was found to be negative for all λ Bootis stars by Maitzen & Pavlovski (1989, A&AS, 81, 335).

A preselection of stars is therefore possible via classical photometric diagrams in order to perform intermediate resolution spectroscopy of candidates.

Baschek et al. (1984) were the first who used low and high resolution IUE-spectra to distinguish λ Bootis from normal type stars. They concluded that carbon (and possibly nitrogen) might not share the underabundances of the heavier elements. The two main criteria to identify true λ Bootis stars, the strong depressions at 1600Å and 3040Å were also detected in normal type, FHB and cool DA white dwarfs (Jaschek et al., 1985, A&A 152, 439). Holweger et al. (1994, A&A 290, L21) later explained the 1600Å feature: it is caused by quasimolecular absorption leading to a satellite in the Lyman α profile due to perturbations by neutral hydrogen.

As a further step, Baschek & Slettebak (1988, A&A, 207, 112) made the first qualitative UV abundance analysis of ten suspected λ Bootis stars. The result was that the light elements (C, N and O) are slightly overabundant whereas the heavier elements (Mg, Al, Si, Mn, Fe and Ni) are moderately deficient (≈0.5 dex). It should be noted that no clear conclusion for sulphur was derived.

As a summary of all previous efforts, Faraggiana et al. (1990, A&A, 235, 311) tried to establish criteria for the definition of a homogeneous group of λ Bootis stars (only the relevant characteristics deduced in the UV region are listed):

The ratio (C I/Al II) should be greater than 1.5

The bump at 1600Å should be detectable

The discrepancy between the MK and UV type should be greater than two spectral subclasses

With the creation of the IUE Final Archive in 1998 it was possible to access all IUE images which are uniformly processed and calibrated, thus making a homogeneous analysis possible.

Solano & Paunzen (1998, A&A, 331, 633; 1999, A&A, 348, 825) presented the results from a careful analysis of the IUE final archive using the low and high resolution spectra separately. Their goal was to establish criteria in the UV region without using data from the optical domain and to unambiguously distinguish λ Bootis from FHB stars.

From the low resolution spectra they concluded that a star is considered to be a member of the λ Bootis group if:

(CI 1657/Al II 1670) > 2

(CI 1657/Ni II 1741) > 2.5

These limits have been set in such a way that they are valid for the entire range of effective temperatures. Small uncertainties due to variations in log g and log Z have also been taken into account to make these limits independent of physical parameters such as the effective temperature and surface gravity. Moreover, they are also independent of the evolutionary status since a clear distinction between λ Bootis and FHB stars was established. From the analysis of the high resolution spectra they were able to add the following criteria for a star being a member of the λ Bootis group:

(C I 1657/Si II 1527) > 3

(C I 1657/Ca II 1839) > 8

The four listed criteria in the UV region are unambiguous and do reflect the typical abundance pattern for this group.

Three papers by Paunzen et al. (1997, A&AS, 124, 23; 1998, A&A, 335, 533; 2002, A&A, 392, 515), summarize the overall pulsational characteristics of the group members (with only two so far not photometrically investigated). They have analyzed the pulsational stability of the individual objects with typical upper limits for bona-fide "constant" objects of about 2 mmag (using also the photometric data of the Hipparcos satellite).

From all well established members (not all within the classical instability strip), 33 have been found to show typical δ Scuti type pulsation. Note that the pulsating spectroscopic binary systems HD 66491 and HD 111786 were not included since their λ Bootis type nature is not well established by now. With the already published abundance values, they have also searched for the existence of a Period-Luminosity-Color-Metallicty relation. They compared the characteristics of the λ Bootis group to a sample of "normal" δ Scuti pulsators. The latter was chosen such that it matches the λ Bootis stars within the global astrophysical parameters. The following properties of the λ Bootis stars are different from those of the δ Scuti pulsators:

At least 70% of all λ Bootis type stars inside the classical instability strip pulsate (this number for "normal" stars is much less)

Only a maximum of two stars may pulsate in the fundamental mode but there is a high percentage with Q < 0.020 d (high overtone modes). For δ Scuti stars this is just the opposite.

The instability strip of the λ Bootis stars at the ZAMS is 25 mmag bluer in (b - y)₀ than that of the δ Scuti stars.

No clear evidence for a log Z correlation with period, luminosity or color was found. The Period-Luminosity-Color relation is the same as for the classical δ Scuti stars.

Bohlender et al. (1999, A&A, 350, 553) have performed high resolution, high signal-to-noise spectroscopy (spectra centered at NaD 5892Å and/or 4500Å) of λ Bootis stars in order to investigate the incidence of high-degree nonradial pulsation. The discovered spectrum variability is very similar to that seen in rapidly rotating δ Scuti stars (Kenelly et al., 1992, ApJ, 400, L71). It has to be noticed that for all but one of the investigated pulsators, high-degree nonradial modes were detected spectroscopically (Bohlender et al. 1999).

Photometric as well as spectroscopic variability is therefore a common phenomenon in λ Bootis stars. Several λ Bootis stars are therefore excellent candidates for multisite photometric and spectroscopic campaigns in order to put further constraints on stellar interior and pulsation models.

Up to now, five λ Bootis stars were targets of multisite campaigns: HD 15165, HD 105759, HD 111786, HD 142994 and HD 192640 (Desikachary et al. 1979, MNRAS, 188, 67; Martinez et al. 1998, MNRAS, 301, 1099; Paunzen et al. 1998, A&A, 330, 605; Mkrtichian et al., 1998, Contributions of the Astronomical Observatory Skalnate Pleso, Vol. 27, No. 3, p. 476) For none of these objects a mode identification was up to now possible.

The photometric data, stellar parameters and calibrated values were taken from Paunzen et al. (2002, MNRAS, 336, 1030). They are based on several calibrations within the Johnson UBV, Strömgren uvbyβ and the Geneva 7-color system as well as data from the Hipparcos and Tycho database. These tables summarize all relevant information about basic properties and astrophysical parameters of this group.

From the Hertzsprung-Russell diagram for all well established λ Bootis stars one is able to conclude that this group consists of true Population I type objects which can be found in the area of the whole main sequence. Paunzen et al. (2002) concluded that the age distribution has a peak at a rather evolved stage (≈1 Gyr) which is in line with the distribution of "normal" A-type stars in the solar neighbourhood.

Using Ti and Ca as examples, Michaud & Charland (1986) showed that this model is capable of producing metal underabundances with the ad hoc assumption of a mass-loss rate of 10^-13M_sun/yr. This high mass-loss rate impedes He settling enough to prevent the disappearance of the superficial He convection zone. Material originally located deep in the envelope has then time to be advected to the surface. Michaud & Charland (1986) showed that the required underabundances materialize after 1 Gyr at the end of the MS lifetime for a normal A-type star.

It was often quoted that this model only predicts underabundances of a factor five which is much too low compared to observations. But considering the poorly known parameters such as the diffusion coefficients, the scale heights and the boundaries of the convection zone, it might well be possible to obtain larger underabundances with this model.

The real problem is the rotational velocity and therefore meridional circulation. What manifests in a cut-off rotational velocity for AmFm stars (≈90 km/s) also destroys the predicted λ Bootis pattern. Charbonneau (1993, ApJ 405, 720) showed that even moderate equatorial rotational velocity of 50 km/s prevent at any time of the MS evolution the appearance of the underabundance pattern because of mixing in the stellar atmosphere.

This very restrictive result seems to reject the diffusion/mass-loss model as a possible explanation for the λ Bootis phenomenon. But one has to keep in mind the simplifications of the used model such as: 1) no evolutionary effects were considered; 2) No turbulence theory was applied and 3) only two elements (Ti and Ca) were considered. Taking all this into account, a definite rejection of this model seems to be premature.

The accretion/diffusion model: Venn & Lambert (1990, ApJ, 363, 234) were the first who noticed the similarity between the abundance pattern of λ Bootis stars and the depletion pattern of the interstellar medium (ISM) and suggested the accretion of interstellar or circumstellar gas to explain the λ Bootis stars. In the ISM metals are underabundant because of their incorporation in dust grains or ice mantles around the dust grains. The depletion factors for various elements can be found in Jenkins (1988, In: Proceedings of the Celebratory Symposium on a Decade of UV Astronomy with the IUE Satellite, ESA, p. 407). Assuming a total particle density of n_tot= 3 cm^-3, the logarithmic depletions D(m) of various metals m range from -0.13 (N) and -0.37 (O) to -0.71 (Mg), -2.02 (Fe) and -2.68 (Ti). Zn shows only a moderate depletion of -0.49. Venn & Lambert (1990) proposed a simple model to explain the depletion variety seen in the atmospheres of λ Bootis stars: the atmosphere consists of a mix of stellar (i.e. solar) and interstellar gas in the proportions (1 - f) to f. The abundance ε of a metal m in the photosphere is then given by

ε(m)/ε(m)_sun = (1 - f) + f D(m)

Following this, Waters et al. (1992, A&A, 262, L37) worked out a scenario, where the λ Bootis star accretes metal-depleted gas from a surrounding disk. In this model, the dust grains are blown away by radiation pressure and coupling between dust and gas is negligible. Considering the spectral-type of λ Bootis stars, the gas in the disk remains neutral and hence does not experience significant direct radiation pressure. The authors showed, that these conditions hold for mass accretion rates below 10^-8M_sun/yr assuming that the gas-to-dust mass ratio in the disk is 100 and that the disk consists of 0.1 μm carbon grains.

Turcotte & Charbonneau (1993, ApJ, 413, 376) calculated the abundance evolution in the outer layers of a 10^-15 and 10^-12M_sun/yr. Solving the diffusion equation modified this time by an additional accretion term, they obtained the time evolution of the Ca and Ti abundance stratification with and without stellar rotation. From their calculations, a lower limit of 10^-14M_sun/yr is derived for the accretion/diffusion model to produce a typical λ Bootis abundance pattern. Moreover their rotating models provide evidence that meridional circulation cannot destroy the established accretion pattern for rotational velocities below 125 km/s. Since diffusion wipes out any accretion pattern within 10 Myr a large number of λ Bootis stars should show observational evidence for the presence of circumstellar material.

Turcotte (2002, ApJ, 573, L129) revised this model recently applying new results on the depth of mixing in A-type stars. Richard et al. (2001, ApJ, 558, 377) showed that the peculiar abundance pattern in Am stars (slow rotators) can only be reproduced by models in which the mixed layers extend much deeper than the superficial convection zone ≈10^-6 times the stellar mass. Inserting this into the accretion/diffusion model rises the accretion rate, which is necessary for the accretion pattern to establish, to 10^-11 to 10^-12M_sun/yr depending on the exact mass of the mixed zone and the rotational velocity of the star.

Andrievsky & Paunzen (2000, MNRAS, 313, 547) worked out the accretion from a circumstellar shell in more detail. They reach the conclusion that gas and dust decouple beyond the condensation radius of a λ Bootis star. Assuming that the dust grains form at the condensation radius, grains will grow only to sizes less than 0.01 μm before they are blown out of the shell. Moreover, the shell accretion can only cause a significant alteration of the photospheric abundances if the density in the shell is proportional to r^-2. Otherwise the shell contains too much material inside the condensation radius, which is not depleted due to grain formation.

All the above discussed scenarios imply a constraint on the evolutionary status of the star, because the existence of a disk or a shell has to be explained in the context of stellar evolution. Circumstellar disks are thought to exist during the pre-main-sequence phase of stellar evolution, while a shell can either occur in a very early phase of pre-main-sequence evolution or after a stellar merger.

Kamp & Paunzen (2002, MNRAS, 335, L45) proposed only recently a slightly different accretion scenario for the λ Bootis stars, namely the accretion from a diffuse interstellar cloud. This scenario works at any stage of stellar evolution as soon as the star passes a diffuse interstellar cloud. The interstellar dust grains are blown away by the stellar radiation pressure, while the depleted interstellar gas is accreted onto the star. Typical gas accretion rates are between 10^-14 and 10^-10M_sun/yr depending on the density of the diffuse cloud and the relative velocity between star and cloud.

Binary theories: Both, the diffusion/mass-loss and diffusion/accretion model have many free parameters and leave some observational facts unexplained like for example the fact that only a small percentage of A-type stars shows the λ Bootis phenomenon or the presence of peculiar hydrogen line profiles. Two ad hoc models addressing these shortcomings appeared in the late nineties. Created ``by resemblance'' these new models connect the origin of the λ Bootis phenomenon to binary systems.

First Andrievsky (1997, A&A 321, 838) proposed the assumption that at least some λ Bootis type stars can originate "by merging" as a result of the dynamical evolution of W UMa contact binary systems. These are close eclipsing binary stars with orbital periods less than one day. Spectral types of both components are almost always similar, mainly in the range between F0 and K0. Due to the angular momentum loss, both components of such a close system approach each other, and finally merge in a more massive but single star. Typical time for merging is ill defined in the literature, we found values from 100 to 200 Myr (Van't Veer 1984, A&A, 139, 477) up to 500 Myr (Leonard & Linnell, 1992, AJ, 103, 1928). The most essential point is that the merging occurs before both stars finish their main sequence lifetime. Since the masses of λ Bootis stars can be found in the range between 1.5 and 2.6 M_sun, their progenitors could be W UMa close binary systems with masses of each component between 0.8 and 1.5 M_sun. In a non-conservative case Andrievsky (1997) proposes that some matter could be lost by the system during the merging phase and form a circumstellar shell.

The above described scenario offers a simple observational test: the search for CNO-processed material which should be mixed into the photospheres of λ Bootis stars as a results of the merger process. But investigations show no abundance anomalies for C, N, O and S in many λ Bootis stars. Andrievsky's suggestion is an attempt to bring into line the apparently evolved nature of some λ Bootis type stars with the requirements of the accretion/diffusion model. It can be neither proved nor rejected, but clearly points to the model's limitations: slow selective accretion has a short time-scale, protostellar shells or disks would be swept out by a stellar wind soon after (≈1 Myr) the ignition of core hydrogen burning and/or photoevaporated by nearby O or B-type stars. One interesting consequence is worth to be mentioned: if some λ Bootis type stars are binary star mergers, they could have a common origin with the well-known blue stragglers (Stryker, 1993, PASP, 105, 1081). On the other hand Holweger et al. (1995, AIP Conf. Proc. 327, p. 41) found through a detailed NLTE abundance analysis, that two blue stragglers in M67 do not show the typical λ Bootis abundance pattern. Instead these two stars fit in the diffusion picture of normal A-type stars.

Abundances determined for the same element by more than one author have been averaged. In a few cases, the abundances differ by more than 0.6 dex and these have been discarded. Abundances for spectroscopic binaries which have been treated as single stars have not been taken into account. In addition, Paunzen et al. (1999, A&A, 345, 597) investigated near infrared (NIR) C and O lines and determined non-LTE abundances for these elements in 21 stars. These abundances are regarded separately, because systematic differences between NIR and optical abundance analyses have been found for these elements. Furthermore, abundances of V, Cu, La, Nd and Eu have been derived for one star by Heiter et al. (1998, A&A, 335, 1009). Paunzen et al. (2002, MNRAS, 336, 1030) list the abundances for individual stars and elements.

This figure shows the abundance pattern, which results from calculating for each element the mean of the abundances derived for all stars. Also shown are the largest and smallest occurring abundances and the number of available abundances per element are given. The mean abundance values, the standard deviations, as well as the differences of the mean to the maximum and minimum abundances are listed in this Table. This figure shows two samples of "normal" stars: 33 normal late A and F dwarfs, which are members of the Hyades cluster (Varenne & Monier, 1999, A&A, 351, 247), and a somewhat more inhomogeneous sample of B to F stars in the galactic field (Heiter, 2002, A&A, 381, 959). All abundances are relative to the solar abundances from Grevesse et al. (1996, ASP Conf. Ser. 99, p. 117).

We conclude from the literature (Heiter 2002; Kamp et al. 2001, Paunzen et al. 2002):

The abundances of C, N and O are on the average solar, but considerable under- and overabundances (-0.8 and +0.6 dex) occur as well.

The star-to-star scatter for the element abundances of C, N, O and S is smaller than for most of the heavier elements.

The mean abundance of Na is also solar, but the star-to-star scatter is ±1 dex.

The iron peak elements are more underabundant than C, N, O, and S for each star.

The iron peak elements from Sc to Fe as well as Mg, Si, Ca, Zn, Sr and Ba are depleted by about -1 dex relative to the solar chemical composition.

Al is slightly more depleted (-1.5 dex) and Ni, Y and Zr are slightly less depleted.

The star-to-star scatter is twice as large as for a comparable sample of normal stars.

Ad 3.: The enhanced abundance of Na has first been noted by Stürenburg (1993, A&A, 277, 139), who derived a mean abundance of +0.6 dex for his sample of stars, and by Paunzen et al. (1999, A&A, 351, 981).
Ad 5.: Note that the mean abundance of Zn is as low as that of the iron peak elements, although its condensation temperature is similar to that of S and should prevent the depletion of this element by chemical separation within an accretion scenario. All available Zn abundances have been determined from the two Zn I lines.
Ad 7.: More than half of the 15 elements which are on the average depleted by -0.4 to -1.5 dex have solar abundance in some stars, and Sr and Zr are even partly overabundant. Extreme cases are Na and Mg. The abundance ranges are only comparable to that of the normal stars for Ni, Y and Ba, and for C and Ca in the normal field stars. Additionally, there are differences in the symmetries of the abundance distributions for Na, Si, Sc and Y. Note that the mean abundance of Al is the lowest in the normal field stars (almost -0.5 dex) as well as in the λ Bootis stars (-1.5 dex), and that the abundances of S in the normal field stars reach remarkably high values. It is thus possible that this sample of normal stars is contaminated with mildly chemically peculiar stars.

The results suggest the existence of a separate chemically peculiar group of "λ Bootis stars" with a characteristic abundance pattern. On the other hand, the scatter of abundances for each element indicates that the "λ Bootis group" is rather inhomogeneous. Furthermore, the comparison to normal stars is difficult because the sample of "normal" main sequence stars with known abundances and parameters similar to that of the λ Bootis stars is rather limited.

Another way of characterizing the chemical composition of λ Bootis stars is to examine abundance ratios between different elements. The most important finding was an anticorrelation between the ratios of light elements (C, N, O and S) to heavier elements and the heavy element abundance. This figure shows that [C/Fe]≥0 for all λ Bootis stars and the iron abundances are significantly lower than those found for superficially normal stars. On the other hand, there is a large overlap with normal stars for other heavier elements.

Heiter et al. (2002, A&A, 381, 971) find no or an only poor correlation of individual abundances with astrophysical parameters such as the effective temperature, surface gravity, projected rotational velocity, age and pulsational period. Within the accretion/diffusion model, the abundances depend mainly on the accretion rate, the mass of the convection zone, and the radiative acceleration; the latter two of them depend mainly on the effective temperature. Therefore underabundances are most likely to appear in a certain temperature interval, in which accretion dominates the diffusion processes. This interval begins at 7000 K for all elements, and its width depends on the element and the accretion rate. For Ti, it extends at least to 9500 K, for Mn at least to 8000 K, and for Ca well over 12000 K. For higher temperatures, the abundance effects depend more on the accretion rate than for lower temperatures. This implies that abundance variations from element to element should be more pronounced among the hotter λ Bootis stars than among the cooler ones. Actual calculations have only been carried out for the elements Ca, Ti, Mn and Eu, for a limited set of parameters, but in the following we assume that these predictions extend to all heavy elements.

The effective temperatures of λ Bootis stars can be found in this table. For a simple test we calculated the standard deviations σ[X] of the heavy element abundances (relative to solar) for individual stars, where at least seven of the elements Mg, Si, Ca, Sc, Ti, Cr, Fe, Sr and Ba have been measured. We did not find any dependence of the abundance scatter on effective temperature. For Mn and Eu the temperature interval in which underabundances occur has been estimated to be rather narrow (Charbonneau, 1991, ApJ, 372, L33). The mean Mn abundance of the five λ Bootis stars with temperatures from 6800 to 7250 K is -1.2±0.3 dex, whereas the abundances of the remaining three stars (∼8000 K) span a large range (-0.4, -1.0 and -2.1 dex), which might be caused by different accretion rates. Due to the lack of Eu abundance measurements, the behavior of this element cannot be tested. This discussion shows that the relation between current abundance and temperature data is inconclusive with regard to predictions of the accretion/diffusion theory.

Another approach to test the basic conditions for an accretion scenario s to investigate the similarities between λ Bootis star underabundances and element depletions in the interstellar gas in detail. The chemical composition of the interstellar gas has been studied by several authors along many different sight lines. Savage & Sembach (1996, ARA&A, 34, 279) give an extensive review on abundances in the interstellar medium (ISM). Additional references for ISM abundances and a table of abundances for six elements for all available sight lines are given in Heiter et al. (2002). For Na, only absorption lines of the neutral element have been studied, whereas singly ionized Na is more abundant in the ISM. Therefore these abundances cannot be compared directly to that of the other elements, except for the sight line to ζ Oph. For this star, a ratio of the ionized to neutral Na column density of 3.1 was derived by Morton (1975, ApJ, 197, 85).

This figure shows the ISM abundances averaged over all sight lines as well as the highest and lowest measured relative abundances. For most elements the abundances vary considerably for the different regions that have been analysed. The scatter of element abundances within the two samples is similar, except for O, which seems to be distributed very homogeneously throughout the ISM. Another similarity between the λ Bootis group and the ISM are the mean abundances of C and O, which are slightly below solar. For the abundances of Na, the ratio of ionized to neutral Na in the direction of ζ Oph has been adopted for all other sight lines, although the validity of this assumption is uncertain.

For S, the ISM abundance distribution is exactly the same as for the λ Bootis stars. On the other hand, the mean abundance of Zn is significantly lower (by 0.6 dex) in λ Bootis stars than in the ISM. Also, the mean abundances of Mg, Al, Si and Mn are slightly lower (but note that only four ISM measurements are available for Al). Furthermore, for all of these elements except Al the lowest abundances are lower in the λ Bootis stars than in the ISM by 0.4 to 0.8 dex. This means that for some stars the underabundances exceed the lower limit set by the ISM abundances in the case of accretion of interstellar gas. For all other heavier elements, the mean ISM abundances lie well below the mean λ Bootis abundances, and also the lowest ISM abundances are lower than that of the λ Bootis stars. The most controversial abundances are that of Zn, which has a mean abundance and distribution similar to that of S in the ISM. The mean [Zn/Fe] ratio is +1.0±0.3 dex in the ISM, whereas the mean ratio of the same elements for λ Bootis stars is +0.3±0.4 dex. This problem is not reduced when the variations in density and temperature of the interstellar clouds are taken into account.

King (1994, MNRAS, 269, 209) searched the IRAS catalogues for infrared detections of λ Bootis stars from the RFB catalogue Renson et al. (1990, Bull. Inform. CDS, 38, 137). He found that about 20% of the bona-fide λ Bootis stars (HD 31295, HD 79108, HD 125162, HD 130158, HD 161868, HD 172167) show an excess in at least one IRAS band. Moreover he concluded that the minimum gas mass necessary to cause the peculiar abundance pattern in λ Bootis stars is so low, that the associated dust mass ∼ 10^-7M_sun may even be below the detection limit of IRAS or submillimeter telescopes.

Using ISOPHOT data from the Infrared Space Observatory (ISO) Fajardo-Acosta et al. (1999, ApJ, 520, 215) confirmed the infrared excess of HD 110411 and report a tentative detection of HD 192640 (only at 20 μm).

Paunzen et al. (2003, A&A, 404, 579) performed a systematic search for dust around λ Bootis stars using available literature data on near infrared photometry, the IRAS catalogues and the ISO archive. They report an infrared excess for the following 8 out of 34 λ Bootis stars: HD 11413, HD 31295, HD 74873, HD 110411, HD 125162, HD 210111, HD 221756 and the binary pair HD 198160/1. Based on the new data reduction, they question the detection of an infrared excess around HD 192640. For HD 183324 the stellar photosphere is detected at 20 and 60 μm excluding any excess radiation due to dust at these wavelengths.

Vega, if regarded as a λ Bootis star, is the only star of this group for which the infrared excess has been analyzed in more detail. The disk as seen by IRAS has a FWHM of 30'' (Aumann, 1991, In: The Infrared Spectral Region of Stars, C. Jaschek, & Y. Andrillat (eds.), Cambridge University Press, Cambridge, p. 363) corresponding to a disk radius of ∼115 AU (Heinrichsen et al., 1998, MNRAS, 293, L78) observed Vega with the ISOPHOT instrument on board of the ISO satellite and found that the disk is resolved at 60 and 90 μm yielding a FWHM of 22 and 36'' respectively. The disk is, as the star (Gulliver et al., 1994, ApJ, 429, L81), seen pole-on. Submillimeter images of Vega reveal an offset of 9'' with respect to the stars position and a FWHM of 24x21'' (Holland et al. 1998, Nature, 392, 788). The typical grain size from fitting the infrared excess with a single-size model is 80 μm (Backman & Paresce, 1993, In: Protostars and Planets III, E. Levy, & J. Lunine (eds.), The University of Arizona Press, p. 1253). Recently Ciardi et al. (2001, ApJ, 559, 1147) using the Palomar Testbed Interferometer found tentative evidence for a disk inside 4 AU contributing about 5% of the flux at 2.2 μm.

It is still difficult to answer the question whether circumstellar dust is a basic characteristic of λ Bootis stars or just as frequent or rare as for all other A-type stars. The general problem is that in most cases the detectors lack the sensitivity to detect the stellar atmospheres at infrared wavelength. This makes it difficult to exclude an infrared excess in the case of a non-detection.

Habing et al. (2001, A&A, 365, 545) overcame this problem by choosing bright stars for which even the photosphere was detectable with the ISOPHOT instrument at 60 μm. Hence a pure photospheric detection ruled out any dust infrared excess. In this way they concluded that 6 out of their 15 A-type stars have circumstellar dust and that these are mostly the younger ones between 200 and 400 Myr (post-main-sequence tracks were used to determine the stellar age of the A-type stars). The general conclusion is that most stars arrive at the main-sequence surrounded by a disk which then disappears within about 400 Myr.

One finds several attempts in the literature to link the λ Bootis stars to the class of Vega-type stars or more precisely the dusty A-type stars like β Pictoris. King & Patten (1992, MNRAS, 256, 571) suggested that β Pictoris might be a λ Bootis star on the basis of its evolutionary status, its circumstellar material and its metallicity photometric indices. Later Holweger et al. (1997, A&A, 320, L49) ruled out this hypothesis by a detailed abundance analysis proving the solar composition of β Pictoris.

Dunkin et al. (1997, MNRAS, 286, 604) performed a detailed abundance analysis of four dusty A-type stars and found no abundance pattern similar to that of the λ Bootis stars. They did find underabundances of Si in two stars and Mg in one star which they tentatively link to the accretion scenario proposed by Venn & Lambert (1990) hinting simultaneously at the apparent contradiction arising from the solar abundances of Ca, Fe and Ti.

There are different criteria to distinguish between circumstellar and interstellar origin. The doublet line ratio of Na I for example is a measure of the optical thickness of the absorbing medium. If the lines are optically thin, the ratio is 2.0, while it decreases to 1.0 in the case of fully saturated lines (see e.g. Spitzer, 1968, Diffuse Matter in Space, Interscience Publ., New York). Moreover an equivalent width ratio W(Ca II)/W(Na I) much larger than 1.0 points towards a circumstellar origin (Lagrange-Henri et al., 1990, A&A, 227, L13).

Gray (1988) was the first to notice a weak narrow absorption core in the Ca II K line of HD 111786, which he suspected to be of circumstellar origin.

Holweger & Stürenburg (1991, A&A, 252, 255) report the presence of narrow absorption components in the Na I D and Ca II K lines for three λ Bootis stars, namely HD 111786, HD 193256 and HD 193281. The ratio of the two Na I D lines is found to be ∼2.0 in the case of HD 193256 and HD 193281, a possible indication for their interstellar origin. HD 111786 on the other hand has a ratio of 1.3, hence a larger optical depth. This star turned out to be in fact a spectroscopic binary consisting of a broad-lined λ Bootis star and a narrow-lined F-type star (Faraggiana et al. 1997, A&A, 318, L21).

Bohlender & Walker (1994, MNRAS, 266, 891) report the detection of a circumstellar shell around HD 38545 from narrow absorption features in Fe II, Ti II, Ca II and the Balmer lines.

Hauck et al. (1995, A&AS, 109, 505; 1998, A&AS, 128, 429) found evidence for circumstellar shell features in the Ca II K line for nine candidate λ Bootis stars and they confirm the circumstellar shell found for HD 38545. They use Crutcher's formula (Crutcher, 1982, ApJ, 254, 82) to derive the velocity component of the ISM in the direction of the star and conclude that their features are clearly circumstellar.

Andrillat et al. (1995, A&A, 299, 493) examined a sample of 20 candidate λ Bootis stars in the infrared spectral region and found that seven of them exhibit evidence of circumstellar shells. HD 38545, HD 39283, HD 84948, HD 98353, and HD 217782 show clear spectroscopic signatures of shells as seen in classical A-type shell stars. Since according to the membership criteria, shell stars should be excluded from the group of λ Bootis stars, none of these stars is presently considered to be a λ Bootis star. HD 84948 was later found to be a spectroscopic binary system. The shell evidence for HD 125162 and HD 183324 is still dubious.

Holweger et al. (1995) found narrow absorption cores in Ca II K in five of the eleven observed λ Bootis stars, namely HD 11413, HD 111786, HD 193256, and HD 198160/1. The presence of these features is correlated with stellar properties like gravity and rotational velocity indicating a circumstellar rather than interstellar origin.

Grady et al. (1996, ApJ, 471, 49) confirmed the detection of shell lines in the spectrum of HD 38545 and announced further the first detection of accreting circumstellar gas in a λ Bootis star. They observed redshifted narrow absorption components of Fe II and Zn II in the mid-ultraviolet resembling those regularly observed in the spectrum of β Pictoris. Following the currently used criteria for membership, this star is no longer a λ Bootis star.

Holweger et al. (1999, A&A, 350, 603) extended the earlier survey for gas around λ Bootis stars (Holweger & Rentzsch-Holm, 1995, A&A, 303, 819) to eighteen stars, and found one more, HD 38545, to show evidence for narrow absorption components in Ca II K. They argue that the Ca K feature seen in HD 111786 can not be produced by the binary nature of this star, that is by the superposition of a λ Bootis star and an early F star. Moreover the strong variability of this feature seems to eliminate an interstellar origin. HD 111786 is not on the master list of λ Bootis stars because of its binary nature. For HD 111786, the equivalent width ratio W(Ca II)/W(Na I) could be determined and was found to be about unity. Since this occurs also for stars where a circumstellar origin has been proven, this equivalent width ratio could also indicate a range of different ionisation conditions or elemental abundances in the gas phase.

Bohlender et al. (1999, A&A, 350, 553) reported the possible detection of strong circumstellar Na I D features in the spectra of HD 142994 and HD 221756. Moreover they note the presence of a most likely interstellar feature in HD 192640. They do not confirm the detection of circumstellar gas in HD 183324 and cautioned against a circumstellar origin of the narrow absorption features in HD 193256 and HD 193281.

Kamp et al. (2002) noted the variability of the narrow absorption feature in HD 38545 as well as an equivalent width ratio W(Ca II)/W(Na I) below unity. The variability underlines the presence of circumstellar gas.

The "Binary star" connection of λ Bootis type stars can be traced out not only between over-contact binary mergers and metallic line weakness which are both discussed by theories, it also offers the ultimate way to shed light on the evolutionary status of these objects. In their extreme form the two leading theoretical models suggest two quite different scenarios for the λ Bootis type stars. Diffusion/mass-loss mechanism will produce the observed abundance anomalies not earlier than at the end of the stars main sequence evolution. Just in opposite, some of the accretion/diffusion scenarios require stars to be very young and even in the pre-main sequence phase of their evolution. To decide about the evolutionary status one needs to compare the position of the star in the Hertzsprung-Russell diagram with the theoretical evolutionary tracks. A knowledge of the stellar mass is crucial for an unambiguous age determination, but deriving the mass for a single star is not an easy task. Binary stars present here a helpful way out.

Two spectroscopic binary systems with metal-weak components are known until now (Abt, 1984, ApJ, 285, 247 classified HD 84948 and HD 171948, both as λ Bootis). HD 84948 comprises two components with almost equal effective temperatures (6600 K and 6800 K), while the absolute bolometric magnitudes are rather different: 0.38 and 1.75, respectively (Iliev et al., 2002, A&A, 381, 914). HD 171948 consists of two quite similar components, both of them have an effective temperature of 9000 K, and an absolute bolometric magnitude of 1.76. The true λ Bootis nature of all four components of these binaries is well-established (Paunzen et al. 1998, A&A, 329, 155). The light elements were found to be moderately underabundant or nearly solar abundant (carbon: -0.8 dex and -1.2 dex for HD 84948 (A and B) and for HD 171948 (A and B), oxygen: -0.6 dex and +0.2 dex for HD 84948 (A and B) as well as -0.4 dex for HD 171948 (A and B)). The underabundances of heavy elements are typically larger than 1 dex and the λ Bootis abundance pattern is very well expressed especially in HD 171948. The galactic space motion of both systems is also typical for Population I stars.

Both spectroscopic binary systems show no eclipses and their orbital solutions contain only the mass-ratio of the two components. Theoretical main sequence and pre-main sequence evolutionary models predict significantly different mass-ratios, making it thus possible to distinguish between the two evolutionary phases. The distinction is easier far from the ZAMS line, where the separation between both types of evolutionary tracks (for a star with given mass) is larger. The location of HD 171948 in the Hertzsprung-Russell diagram is very close to the ZAMS, and both evolutionary models predict a mass-ratio of about 1.0. The case of HD 84948 is more promising: both components of this binary system are far from the ZAMS, and pre-main sequence models predict a mass-ratio of about 1.5, while main sequence models give a value of 1.28. More than sixty high signal-to-noise spectroscopic observations of both systems allowed (Iliev et al. 2002) to determine the orbital elements and the mass-ratios. For HD 84948 they got a mass-ratio of 1.17, ruling out its pre-main sequence nature at a 10σ level. The mass-ratio measured for HD 171948 is close to the predicted one of 1.0.

Evolutionary ages could be determined along with the masses from the stellar evolutionary models. Using computations of Claret (1995, A&AS 109, 441) an age of about 1 Gyr for HD 84948 was obtained. For the first time direct evidence was found that at least some bona-fide λ Bootis type stars are near to the end of their main sequence lifetime. With its age of about 10 to 100 Myr HD 171948 is close to the ZAMS. It is older than other λ Bootis type stars found in the Orion OBI association and NGC 2264, but it is substantially less evolved than HD 84948. Thus, the study of the two spectroscopic binary systems HD 84948 and HD 171948 provides an indication that the λ Bootis phenomenon takes place at any stage of main sequence evolution.The diffusion/mass-loss and the diffusion/accretion theories, which rely on a pre-main-sequence evolutionary state, seem to be inadequate to explain the observed abundance anomalies.

The λ Bootis stars in the UV region

A statistical investigation of the spectroscopic binary frequency among λ Bootis stars

Na-D observations of λ Bootis stars