You are an expert at summarizing long articles. Proceed to summarize the following text: 
to characterize the metallicity and age of an old stellar system , iron abundance [ fe / h ]   ( the logarithm of the iron - to - hydrogen abundance ratio with respect to its solar value ) and the stellar effective temperature  must be determined together for its main - sequence turnoff ( msto ) stars . these are dwarfs and subgiants of types f and early g with @xmath0 . their spectra are dominated by fei  lines , whose strength increases as  declines . for both metallicity and age , then , a reliable derivation of  is necessary . this can be problematical even for an individual star . the photometric colors often used to derive  are sensitive to reddening by the interstellar medium and to the modeling of convection ( castelli , gratton , & kurucz 1997 ) , as is the more recently developed infrared flux method @xcite . methods using balmer - line profiles are reddening - independent but still very model - dependent @xcite . methods relying on the excitation equilibrium of fei  are reddening - independent and model - independent if lines are weak , but are susceptible to systematic errors in laboratory measurements of gf - values @xcite and to possible departures from local thermodynamic equilibrium ( lte ; * ? ? ? * ) . currently ,  values derived for individual msto stars show uncomfortably large spreads : values of @xcite and @xcite average 100 200  lower than those found by @xcite and @xcite , with those of @xcite and @xcite in between . not only is overall metallicity affected , but also the ratios of the abundance with respect to iron of certain elements . among these are the  and  ratios so critical to uncovering the nucleosynthesis history of the early universe @xcite and deducing the age of globular clusters from the msto luminosity @xcite . when  alone is raised by 100 , modeled values of   change by + 0.09dex @xcite ;  changes by @xmath1dex , + 0.03dex , or + 0.16dex depending on whether the high - excitation oi  7775  triplet , the [ oi ] 6300  ground - state line , or oh molecular lines near 3100   are used @xcite . not surprisingly , the oxygen abundance in extremely metal - poor stars remains uncertain , as summarized during a workshop devoted to the problem @xcite . @xcite and @xcite , among others , have suggested that a higher  scale for metal - poor turnoff stars might eliminate the systematic differences in  seen from its different diagnostics .    for all these reasons , it seems important to reduce uncertainties in the determination of  in msto stars . fitting the slope of their uv flux distribution might help , since the fluxes of f g stars peak in the optical so that uv fluxes vary rapidly with . although fitting a uv slope is also reddening - dependent , reddening should be very low for the brightest metal - poor msto stars , which are typically within 50pc ( e.g. , * ? ? ? reliable calculations of msto mid - uv spectra would be especially valuable in constraining the age and metallicity of a spatially unresolved old stellar system such as an elliptical galaxy or extragalactic globular cluster . for the mid - uv flux of an old stellar population is dominated by turnoff stars @xcite , but the light redward of 4500  is dominated by cool giants in a 17  population of solar metallicity @xcite . to date , however , mid - uv studies of galaxies and globulars ( e.g. , * ? ? ? * ; * ? ? ? * ) have largely been empirical , comparing extragalactic spectra against observed spectra of individual f and g stars @xcite . thus the metallicity of the extragalactic systems has generally been assumed to be solar , although age is very sensitive to this assumption @xcite . furthermore , nonsolar abundance ratios have not been included , despite the fact that abundance enhancements of light elements such as magnesium are generally seen in metal - rich extragalactic systems @xcite . modeling the mid - uv region is difficult , though , because of the dramatic increase in line absorption at short wavelengths . line blending in the solar spectrum suppresses the optical continuum blueward of @xmath2 4500   @xcite , and becomes even more severe in the near- and mid - uv . especially in the uv , many lines are `` missing ''  although they appear as absorption features in spectra of solar - type stars , they lack laboratory identification , and are thus not matched by spectral synthesis calculations @xcite . as a result , the normalization of the solar uv continuum is controversial @xcite .    with few exceptions ( e.g. , * ? * ) , calculations of the mid - uv region are based on kurucz model flux distributions . @xcite compared these to international ultraviolet explorer ( iue ) mid - uv spectra of nearby metal - poor halo stars , illustrating that  is generally recovered when parallaxes are known and reddening negligible . @xcite have used kurucz fluxes as generated by @xcite to examine the effect of abundance on the mid - uv light of old stellar populations , but found that the uv fluxes of cool stars were poorly reproduced . this is not surprising given the known limitations of the opacity distribution functions on which the kurucz flux distributions are based . they statistically treat lines in a list that includes `` predicted '' lines whose wavelengths are uncertain by @xmath210  ( because one or both energy levels have not been measured in the laboratory ) , limiting the resolution of the flux distributions to 10 . moreover , they assume a solar iron abundance @xmath20.15 dex higher than the value currently accepted @xcite , and a solar ratio of light - element abundances  = 0 . in contrast , metal - poor stars of the galactic halo typically show  = 0.3 @xcite , a factor - of - two enhancement in mg , si , ca , and ti . this affects the fit in the mid - uv , since it includes the very strong resonance doublet of mgii  at 2800   and a bound - free opacity edge of mgi  at 2512 , plus many weaker mg , si , and ti lines . in this work we calculate the mid - uv spectrum line - by - line at high resolution , independent of the kurucz flux distributions . our line list does not include predicted lines , but rather is based on an updated laboratory line list augmented by manual addition of missing lines where necessary . to ensure the reliability of these calculations , we match as best we can the individual line blends observed in the mid - uv echelle spectra of real near - turnoff stars , spanning a wide metallicity range and including both solar and enhanced light - element ratios . our purpose is to reproduce mid - uv spectra of turnoff stars accurately enough to determine  and to predict the composite mid - uv spectra of old stellar systems over their entire range of overall metallicity and light - element abundance ratio .    in this first step , we describe this calibration procedure . because the calibrations depend significantly on the stellar temperatures adopted , we have carried out new  determinations in seven metal - poor turnoff stars and one mildly metal - poor hotter star . they are based on simultaneously matching both mid - uv spectra ( 2280  3120 ) and the  profile and the strengths of high - excitation atomic lines in optical spectra . we thus do not rely on any photometry or parallax information in the  determination . iron abundances and  ratios were adopted as derived from both weak and strong lines in the optical spectra . to assist in filling in the mid - uv line list , we have included mid - uv analyses of two stars with mid - uv spectra of very high resolution , the sun and 128620 (= ) , a well - studied nearby southern star with a greater - than - solar metallicity @xcite , for which  was taken from the literature .    in fitting each mid - uv spectrum , we adopted a single continuum normalization constant , the scale factor by which the observed mid - uv spectral flux was multiplied to match the flux predicted by theoretical calculation . this constant proved to agree to an average of 6% @xmath3 5% with that expected from the ratio of stellar distance ( found from the parallax ) to stellar radius ( found from the model  assuming a reasonable stellar mass ) . model , , and   colors all agreed extremely well with those observed , confirming  errors of @xmath450 . our mid - uv and optical echelle spectra are described in  3 . as discussed in  4 , was initially established from the optical spectra , and compared to the mid - uv fit . for the most metal - poor stars , this fit did not require the addition of missing lines , but did require the use of kurucz models in which convective overshoot was turned off , in contrast to the kurucz models in general use . to reproduce mid - uv line strengths and to match the pseudocontinuum of the mid - uv spectra of more metal - rich stars , mid - uv atomic line parameters were changed and missing lines were added . transitions especially sensitive to the surface temperature were still poorly reproduced ; this mismatch was reduced by adopting a mild enhancement of temperatures in shallow photospheric layers of the models of one - tenth solar metallicity and higher .    the degree to which the resulting calculations match both mid - uv and optical spectra of turnoff stars is shown and discussed in  5 . the match is very good for 2900  3120  at all metallicities , and generally good throughout the mid - uv for metallicities at and below one - third solar . it could be improved at solar metallicities by obtaining additional high - resolution mid - uv spectra of a few well - chosen standards . our  results are presented in  6 , where comparisons are made of model fluxes and colors with observations . we conclude in  7 with a summary and a discussion of the prospects for constraining age and metallicity of quiescent stellar systems by comparing such calculations to spectra of their integrated light . we begin by listing the properties of the stars considered in table 1 . the first five columns following the hd number give the observed @xmath5 magnitude , colors , and the parallax from simbad . the next five columns , discussed in  6 , show comparisons of observed properties of each star with those calculated from the models . the parameters of the model used to calculate both mid - uv and optical spectra of each star are given in the next four columns . in all cases but 128620 =  and the sun , all the model parameters were redetermined here as described below . the final columns show  values determined previously in several works . of the eight stars whose temperatures have been redetermined here from optical and uv spectra , seven are metal - poor main - sequence turnoff ( msto ) stars . the eighth , 128167 , is a mildly - metal poor star lying above the halo main - sequence turnoff , either younger or a blue straggler . it was included to give temperature leverage to the mid - uv comparisons , as discussed below . likewise the sun and  were included to represent the high - metallicity end . the solar parameters were considered so well determined that no re - evaluation was necessary . the star  was not reanalyzed because no optical spectrum was available : it is inaccessible from the northern hemisphere . its  and abundance have been determined to moderate accuracy from extensive abundance analyses , and its gravity as  = 4.33 @xmath3 0.02 as found from its visual and radial - velocity motion about its companion @xcite . the iron abundance listed in table 1 was adopted to best match its mid - uv spectrum . optical spectra were obtained for all stars in table 1 but the sun and . the lick observatory shane 3 m echelle spectrograph @xcite was used with a tek @xmath6 ccd , at a fwhm resolution of 38,000 . the spectra of 84937 , 94028 , and 106516 were kindly provided by @xcite . exposure times on the 3 m were typically 10min for all stars except 128167 ( 40s ) . complete spectral coverage over 3900 7800  and s / n @xmath7 100/pixel were obtained in all cases but one , although fringing compromises the accuracy of the data at the long - wavelength end of the range . for the one exception , 114762 , the blue spectrum was obtained in two one - hour exposures by tony misch at lick with the coude auxiliary telescope , using the same echelle spectrograph configuration . for the sun , the flux spectrum of @xcite was downloaded from the kurucz web site at http://cfaku5.harvard.edu ; the full - resolution version was matched at 150,000 resolution . the mid - uv spectral observations of the sun modeled here are the rocket spectra of @xcite , recorded on film . both passes of the center of the solar disk were used . gaussian broadenings of 1.5  for macroturbulence and 80,000 fwhm for other broadening sources provided a good match . a constant normalization scale factor of 1.10 was adopted , in keeping with the stated 10% uncertainty of its flux calibration . the stellar uv spectra were obtained with the space telescope imaging spectrograph ( stis ) of hubble space telescope ( hst ) from mid-1998 to mid-1999 . salient features are noted here ; full details may be found from the hst archive at http://archive.stsci.edu .    to obtain the mid - uv spectra of most stars , the near - uv mama detector was used with the stis e230 m grating over its default range of 2280 3120  at a fwhm resolution @xmath225,000 . 106516 (= hr4657 ) was observed by @xcite in program # 7433 with the @xmath8 aperture . mid - uv spectra for 19445 , 84937 , 94028 , 114762 , 184499 , and 201891 were acquired in snapshot program # 7402 by @xcite , with two exposures obtained for 19445 , 84937 , and 114762 . the @xmath9 aperture was used and exposure times were 6  10min . consequently , s / n is lower , but flux levels are less dependent on pointing . fluxes good to 1% for these six stars are indicated by previous observations of mgii  line profiles in each star , made using the hst goddard high - resolution spectrograph ( ghrs ) and the @xmath10 aperture in program # 5869 @xcite . mid - uv spectra of the stars 128167 ( = hr5447 ) , from # 7433 , and 128620 (= hr5459 =) , from # 7263 @xcite , were recorded in multiple exposures with the e230h grating . spectral coverage is 2380 3160   for 128167 and @xmath112280   3160  for 128620 . a fwhm resolution of @xmath260,000 was achieved with the @xmath12 aperture for 128167 , and the @xmath13nd aperture for 128620 . for neither star do fluxes always reproduce exactly in orders obtained in more than one exposure ; differences as large as 10% are seen . similar uncertainties are anticipated for 106516 , where the @xmath12 aperture was also used . rotation is detected in at least two stars and possibly in two more , from the extra broadening required to match line profiles . for 128167 , 106516 , 84937 , and 114762 , line breadths were matched in the spectral synthesis by adopting rotational velocities  = 9 , 7 , 5 , and 4  respectively . thus the resolution of the mid - uv spectrum of 128167 is effectively degraded to that of the stars observed at lower resolution , affecting the discernment of mid - uv lines . this star nonetheless has one of the sharpest - lined spectra among those obtained in program # 7433 . spectrum analysis was based on visually matching each observed spectrum to _ ab initio _ radiative - transfer calculations assuming static models in lte . the program synthe @xcite , the @xcite and @xcite grids of atlas9 models , and the kurucz line lists for atomic species and molecular hydrides were downloaded from the kurucz web site , where further details on each of these can be found . as described below , all models , including that of the sun , were ultimately drawn from the @xcite grid alone . the more metal - rich ones were modified by introducing enhanced temperatures in the shallow layers . the synthe program accepts as input both line lists and a model atmosphere with temperature and other quantities tabulated for each depth , and calculates an array of spectra at individual emergent angles . these may then be used individually ( as for spectra of the center of the solar disk ) or convolved with a zero or positive rotational velocity to generate an emergent flux spectrum . in either case , the resulting spectrum can be broadened by gaussians ( or other profile shapes ) to incorporate macroturbulence and instrumental broadening . synthe was modified to run under unix on a sun ultra-30 at the university of virginia . mid - uv spectra from 2280  to 3160  were calculated at a resolution of 330,000 ( 500,000 for the sun and 128620 ) , and optical spectra at a resolution of 500,000 . a gaussian macroturbulent broadening of 1.5  was assumed for all stars as well as for the sun . spectra for some stars were rotationally broadened as listed in the last paragraph of  3 , as was the solar flux spectrum , for which a mean rotational velocity of 2  was adopted . the spectra were then broadened by gaussians corresponding to the values of instrumental resolution listed individually in  3 .    the line lists and the source code of the opacity routines available at the kurucz web site give specifics on the wide variety of opacity sources treated by synthe . for example , the significant absorption by the bound - free absorption edge of mgi  near 2512  is treated both as line opacity , with the line list including the bound - bound transitions to @xmath14 , and as continuous opacity blueward of the bound - free edge at @xmath15 = 39759.8 @xmath16 , by the subroutine mg1op . near @xmath15 , this routine adopts a cross section of 40 mbarn with a @xmath17 frequency dependence , based on the laboratory measurement of 46 @xmath3 12 mbarn by @xcite and the convergence of high - level lines . where this formalism falls below the theoretical cross section , the program adopts the latter , with a cross section of 20 mbarn and a @xmath18 frequency dependence ( e.g. , * ? ? ? * ) .    for calculations redward of 5000 , the line list of @xcite based on the solar spectrum was adopted , with modifications due to the use here of a flux , rather than central intensity , spectrum for the sun and to the incorporation of an approximate chromosphere in the solar model . this led to the deduction of larger gf - values for the strongest lines except where laboratory gf - values were adopted , as for low - excitation fei  lines with furnace measurements @xcite . blueward of 5000 , the kurucz 25 may 1998 atomic line lists _ gf_0300.100 , _ gf_0400.100 , and _ gf_00500.100 were used . these differ significantly from those used in the kurucz flux distribution calculations in including only atomic lines identified in the laboratory ( with laboratory gf - values where available )  like those of @xcite but with the addition of many lines of fei , feii , and other species whose upper energy levels were newly determined in near - uv laboratory spectra ( e.g. , * ? ? ? * ) . because the strengths or profiles of many lines were mismatched when using line parameters directly from this list , their gf - values and damping constants were changed where necessary , as described shortly . since many lines appear in the stis mid - uv spectra that were not modeled by lines in the list , such missing lines were added blueward of 3120 . redward of this , no missing lines were added , but gf - values were changed to match stronger lines in the sun in the regions depicted in fig . 4 below . the , , [ fe / h ] , , and  values for each reanalyzed star were initially established from  and weak lines of neutral species in the optical spectra , reassessed from optical lines of once - ionized species , then confirmed or revised from the uv . for the two most metal - poor stars , mid - uv spectra provided an immediate check on . for the others , the following steps were iterated until all diagnostics converged . first the  profile and an assumed  were used to estimate . [ fe / h ]  was adjusted until weak fei  lines in the 5100 5400  region were well matched ; their strengths are largely gravity - independent . the strength of weak tii , mgi , and sii  lines determined . microturbulent velocity  was set by matching weak and strong iron and titanium lines . gravity  was then checked by demanding matches for mgi , mgii , and the wings of the  lines , for fei  and feii , for tii  and tiii , and for sii   and siii . this also constrains , for the mgii  and siii  doublets arise from highly excited lower levels ( @xmath78ev ) , while the optical tiii  and feii   lines originate from much lower levels ( @xmath113ev ) . roughly two dozen fei   lines , a dozen tii  lines , and a dozen total tiii  and feii  lines were detected for the most metal - poor stars , with double this number typical of the more metal - rich stars . the mgii  and siii  doublets at 4481  and at 6347.1 and 6371.3  were seen in all cases , but siii  only marginally so in 84937 and 19445 . once the diagnostics began to converge , iterative adjustments were made to the kurucz mid - uv atomic - line parameters . first , gf - values were modified for atomic lines identified in the laboratory . the relative contribution of identified lines to a blend was ascertained from the behavior of the blend in the hot star 128167 versus the cooler ones , and in the higher - gravity stars such as 19445 versus the lower - gravity ones . damping constants were changed in some cases to match line profiles . damping constants rather than gf - values were changed for the strong resonance lines well - observed in the interstellar medium for which gf - values are tabulated by @xcite . next , for lines lacking identification , the laboratory line list was searched for transitions close enough in wavelength to perhaps be responsible . the gf - value of a nearby identified line was increased by up to 4.5 in the log , and the calculation repeated . this line was dropped if still not strong enough ; otherwise its gf - value was adjusted to match . the high resolution of the  and solar spectra showed that most such identifications were very well matched in wavelength , and thus likely to be correct , when based on increases of @xmath19 2.5 in the log . for the two most metal - poor stars , these steps were sufficient to enable the continuum to be defined unambiguously throughout most of the mid - uv region . its definition in more metal - rich stars was hampered by line crowding , yielding a pseudocontinuum that the modeling did not properly reproduce . much of this pseudocontinuum mismatch appeared to be due to individual transitions missing from the line list . this was indicated by the very common occurrence throughout the mid - uv of unmodeled absorption features found at the same wavelengths in all stars , features whose strength steadily increased in going from weak - lined to strong - lined spectra . most of these transitions are weak , consistent with the fact that the intrinsically strongest transitions of each species are the ones most likely to have been detected in the laboratory .    except just redward of the mgi  edge at 2512 , most missing lines are probably due to iron . its abundance is high and the spectra are rich in lines of fei  and feii , which are well - populated species in solar - temperature stars . of the elements with similar first ionization potentials , only magnesium and silicon have comparable abundances ; the mid - uv spectra of their low - ionization stages are generally much less complex , as are those of the more abundant cno elements . while iron - peak elements such as cr , mn , and ni have many transitions throughout the mid - uv , iron dominates because it is twenty or more times as abundant . in most halo stars above one - tenth solar metallicity , the relative abundances of these elements with respect to iron are within 0.2 dex of those of the sun . the mismatch that will result from assigning missing lines of these elements to iron should be mild .    in the region immediately redward of the bound - free mgi  opacity edge at 2512 , the probability is significant that missing lines may be due to magnesium rather than iron . redward of the edge there are a very large number of transitions to known highly excited levels of the mgi  atom , and presumably a very large number of similar transitions to closely - spaced highly excited levels that are as yet unidentified in the laboratory . if many weak transitions redward of 2512  are due to magnesium but have been assigned to iron instead , our calculations will underestimate the absorption in this region in metal - rich alpha - enhanced populations . we see below that this region is poorly modeled in any case once metallicity exceeds one - tenth solar . the modeling then proceeded iteratively by assigning tentative identifications and transition probabilities to each missing feature , then recalculating until an acceptable match was achieved or deemed unattainable with existing data ( as for the regions just redward of 2550  and 2640  in the solar spectrum ) . features were attributed to either low - excitation fei , moderate - excitation fei , or feii , whose line strengths depend differently on stellar temperature and gravity ( pressure ) because of the boltzmann and saha equations . with a  @xmath21000  higher than those of other moderately weak - lined stars , 128167 proved exceptionally useful in making these assignments . the higher resolution of the spectra of the strong - lined stars 128620 =   and the sun was essential in discerning the wavelengths of the many features emerging above one - third solar metallicity . a brief examination was made of the gf - value changes for low - excitation fei   lines , namely those identified in the laboratory with ( air ) wavelengths falling between 2600  and 3000 , and with a lower level of 6000 10000   and an upper level of 40000  50000 . sixty - five such fei  lines have gf- values from the laboratory measurements of @xcite ; 52 were not changed , while 13 were changed by from @xmath20 to @xmath21 dex . nine lines have gf - values from the laboratory measurements of @xcite ; five were unchanged , and four changed by from @xmath22 to @xmath23dex . thus the mid - uv gf - values found in this work generally concur with the mid - uv laboratory measurements . in contrast , fifteen lines have theoretical gf - values from @xcite ; four were not changed , but 11 were changed by from @xmath24 to @xmath25dex , all but one decreased . the theoretical gf - values of @xcite are seen to be subject occasionally to very large errors , usually overestimates . this can be attributed to the finite size of the cray computer used in those calculations , which dictated that the number of levels included be truncated , and so upper levels were preferentially omitted from the model atom . because the sum of the gf - values of transitions arising from a given lower state is normalized to unity , gf - values for the transitions included in the calculation will be overestimated because of the exclusion of others arising from the same level . thus the calculations will overestimate the gf - value by increasing amounts as the fraction decreases of the actual possible upper states that are included in the atomic model . this fraction decreases to the blue , because higher upper levels are involved for a given lower level . indeed , the severity of overestimates appears to increase to the blue : all the above changes made to the @xcite gf - values were @xmath26 or lower below 2650 , but redward of this only one change this negative was required . the overestimates among the @xcite fei  gf - value calculations are thus seen to be due to incompleteness and not to more fundamental difficulties . undoubtedly many of the weaker mid - uv transitions are missing entirely from the predicted list .    for the hydrides , separate treatment was required , as the lines are naturally associated into bands , and the kurucz hydride calculations are now several decades old . for oh , the theoretical calculations of @xcite ( communicated in august 2000 by m. bessell ) were used wherever possible . for ch , the lifbase data ( e.g. , * ? ? * ) were used ( communicated in august 2000 by m. bessell with lines cross - identified ) . since these cover only the strongest bands , calculations must still rely on the kurucz values for the others , notably the 4 - 3 band of oh with many weak but discernible lines in the sun . a comparison of the kurucz oh transition probabilities against those of @xcite showed strong trends , with discrepancies increasing at high rotation number j. to place them on approximately the same scale , the kurucz transition probabilities for the oh and ch lines were decreased by an amount @xmath27 = @xmath28dex . once missing atomic lines were largely filled in , a comparison was feasible of calculated oh and ch line strengths against those of the sun and stars . because solar oh lines were otherwise modeled too strong , all the transition probabilities were reduced by an additional 0.15 dex uniformly . in part because weak oh lines in the sun were reasonably well modeled but strong ones were still too strong , a chromospheric mitigation of the surface temperature drop was invoked as described shortly . as illustrated below in figs . 3@xmath29 and @xmath30 , these changes suffice to reproduce both strong and weak 0 - 0 oh features in stars as well as the sun . the calculations were first run using @xcite models , but the very metal- poor stars showed a  deduced from  as much as 200  higher than that from fitting the mid - uv . that discrepancy was removed by adopting @xcite models , models in which convective overshoot has been turned off . we did not use the alpha - enhanced @xcite models , as they are available for only a single value of the alpha - element abundance enhancement .    like the @xcite models , the @xcite models we used were calculated with  = 2  and a logarithmic solar iron abundance of @xmath31 with respect to the abundance by number of hydrogen plus helium , which amounts to @xmath32 with respect to hydrogen alone . these choices are about 1  higher in  than the @xmath21  typically found for the solar flux spectrum , and about 0.15dex higher in iron abundance than indicated by the meteoric iron abundance and by several recent determinations for the sun ( e.g. , * ? ? ? * ; * ? ? ? thus both sets of models were calculated with higher line blanketing than expected for a solar - type star of the nominal model metallicity , since iron lines dominate the uv and optical spectrum , and the higher  increases the strengths of moderately strong lines of every species . for each star including the sun , a @xcite model was interpolated to the nearest 50  in  ( 25  for the sun ) , 0.1dex in , and 0.05dex in [ fe / h ] . the microturbulent velocity  was changed at all depths to the value indicated for each star in the table and figures below . to minimize the overblanketing effect , our models were interpolated to a nominal grid abundance 0.10  0.15dex lower than that desired . to set abundances , the abundance of the interpolated model was then raised to the desired value and the relative abundance of iron dropped by 0.15 dex . the specific light - element abundance elevations required by the optical observation were adopted , or assumed to be zero for the sun and , and all species equilibria were recalculated . this procedure incorporates the specific  and iron and light - element abundances set for the model in calculating both the ionization and molecular equilibria and the line and continuous opacities for its spectrum . concerning the model structure , the procedure compensates for the effects of overblanketing due to the high choice of iron abundance , but not for the light - element ratio nor for the microturbulence . we anticipate very mild errors in model colors and temperature structure until spectra become dominated by molecular lines or otherwise are considerably stronger - lined than any of these . the comparison of model and observed colors (  6 ) bears out this expectation . for the sun , a model with  = 5775 ,  = 4.4 , [ fe / h ]  = 0 ,  = 0 , and  = 1.0 , was interpolated in the above way . for 128620 =  we adopted a similarly interpolated model with  = 5800 ,  = 4.3 , [ fe / h ]  = + 0.15dex ,  = 0 ,  = 1.0 . while @xcite found  = 5830  and [ fe / h ]  = 0.25dex , most other analyses have derived somewhat lower  and [ fe / h ]  values . at  = 5800  and  = 1.0 , [ fe / h ]  = 0.15dex provides a better match to the mid - uv spectrum of .    even with the use of @xcite models interpolated as above , another major type of mismatch persisted in all strong - lined stars : cores of lines formed near the surface appeared too strong in the calculations , especially in extremely strong mid - uv lines such as those of mgii  and mgi  or feii  and fei . this is definitely not due to wayward gf - values or damping constants , because of the trend seen in the mismatch . profiles calculated for the strongest lines in the most metal - poor spectra generally matched or were somewhat too high in the core , but were somewhat too low in the core in stars of intermediate line strength , and became too low across the the central ten ngstroms or more in the strongest - lined stars . ( see fig . 3@xmath33 below for residual trends of this nature . ) as noted above , a related mismatch was seen within the oh spectrum itself : in any given strong - lined spectrum , notably that of the sun for which the model parameters are extremely well known , weak oh lines were matched at a higher oxygen abundance than strong lines . the possibility discussed in  5 that the solar continuum may be drawn too low would lead to a mismatch of opposite sign . to the extent that the gf - values are reliable , then , this indicates an error in the temperature structure of the model . the most obvious possibility is a chromosphere , which is not incorporated in flux - constant , static models such as those of synthe . the solar chromosphere , for example , is seen in images of the limb and evident from emission reversals in the cores of strong spectral lines such as mgii . it is generally modeled by an increase in temperature in relatively shallow layers of the atmosphere ( e.g. , * ? ? ? * ) . any of several non - radiative heating sources extending from the photosphere through the chromosphere into the corona could be responsible ( e.g. , * ? ? ? * ; * ? ? ? * ; * ? ? ? we consequently changed the temperature of individual points in each interpolated stellar model with [ fe / h ]  @xmath34 , beginning just below the point where the model becomes optically thin at 5000  ( typically at depth number 51 or 52 ) . ( we tried this for the more metal - poor models as well , but the profiles of all the strong lines became much too high near the core , so we left these models as they were . ) to best fit both optical and mid - uv spectra simultaneously , we increased the temperature at model depths 28 through 48 by 10 , 20 , 30 , 40 , 50 , 60 , 70 , 80 , 90 , 100 , 110 , 120 , 130 , 140 , 150 , 150 , 140 , 120 , 90 , 60 , and 30  respectively . this temperature enhancement may be considered an approximate chromosphere , though it is based entirely on empirical fits of one - dimensional spectral line calculations to observed spectral line profiles . it incorporates neither energy transport considerations nor matches to spatially resolved structures , both of which lie beyond the scope of this work . the degree to which our mid - uv and optical spectral calculations match spectra of real stars is illustrated in the figures . the first three show mid - uv spectra , and the last , optical . within a figure , panels are ordered by wavelength . flux ( central intensity for the solar mid - uv spectrum ) is plotted versus wavelength in air . the spectrum calculated for each star ( light line ) is plotted on top of that observed ( heavy line ) .    in figure 1 , the global mid - uv fit is shown for the nine stars with absolute fluxes . all spectra have been gaussian smoothed to a fwhm resolution of 500 ( 550 for 128127 and ) . each spectrum is displaced vertically by @xmath220% from the one below . for each star a single mid - uv continuum normalization constant was adopted , as listed on the left . on the right , the hd number of each star appears in bold . below are found the model - atmosphere parameters adopted for the corresponding calculation :  ( in ) , c wherever the chromosphere was modified as described at the end of  4 , log @xmath35 , [ fe / h ] ,  ( in ) . on the following line appear the values in dex of the light - element abundance enhancements [ element / fe ] ( in dex ) , first that of oxygen , then those for mg , si , ca , and ti . ( the latter four are grouped together wherever they are the same . ) while the fits are generally good , serious discrepancies are noted in  at the bottom of the figure near 2550  2600  and 2640  2710 , regions of relatively little line absorption in the more metal - poor spectra . figure 2 illustrates the effect on mid - uv fits when model parameters are changed somewhat . two of the stars of fig . 1 are depicted : the moderately weak - lined 94028 ( top two spectra ) and the rather strong - lined 184499 ( bottom three ) . star and model identifications and normalization constants are indicated as in figure 1 . the lower plot for each star represents its best fit , a replot of that star s comparison in figure 1 .    comparing the bottom plot to the middle plot shows the effect for 184499 of decreasing  by 0.3dex to the solar  ratio . although minimized by the adoption of a normalization constant 11% higher , the mismatch is still evident . the strengths of the wings of the strong mgii  doublet at 2800  ( and of mgi  at 2852 ) are poorly reproduced , as is the overall pseudocontinuum level near and below the bound - free edge of mgi  at 2512 . similar calculations run for 94028 show a more dramatic effect in the mgii  and mgi  line wings , which are weak enough that little renormalization is possible , but a smaller effect near 2512 , where the h@xmath36 opacity is dominant because of the reduced magnesium abundance .    comparing the plot just above the best fit for each star to the best fit itself shows the degeneracy of the mid - uv spectra with respect to a simultaneous change in ,  and [ fe / h ]  designed to preserve overall line strength and ionization equilibrium . for each plot just above the best fit , was decreased by 100 , [ fe / h ]  and  both concurrently decreased by 0.1dex , and the spectrum renormalized accordingly . this change is very difficult to detect from fitting the mid - uv slope alone at these metallicities and this resolution . the increased metallicity reduces the blue pseudocontinuum relative to the red at the same time that the increased   raises it . the change is marginally detectable in 94028 from the mismatch of low - opacity windows , but not in 184499 because the overall match is less satisfactory . such a change does lead to a reduction in the normalization parameter , by 11% for 94028 and 17% for 184499 . in principle this might be detected by other means , but only if  is properly modeled . for as seen above , taking  = 0 ( as in the kurucz flux distributions ) can largely compensate in stars such as 184499 that show a light - element enhancement . our conclusion , then , is that if the mid - uv normalization constant is treated as a completely free parameter ( as here ) , the mid - uv flux distribution by itself can currently be used to set  only in msto stars with [ fe / h ]  @xmath37 . while the mid - uv was useful in demonstrating the need for models with no convective overshoot and with an approximate chromosphere , our determinations of  for the more metal - rich msto stars have relied on the optical spectra instead . we show in fig . 4 below that the balmer line profiles and the strengths of other high - excitation lines are well matched in the optical spectra of all stars at the same  as the best - fit mid - uv spectrum , and in  5 that model colors match observations as well , so that  is indeed very well established here . however , our goal of  determination for turnoff stars of all metallicities from mid - uv spectra alone is not yet met . we are not able to calculate the low - opacity windows blueward of 2900  accurately enough at near - solar metallicity , as we now demonstrate . figure 3 plots at higher resolution and higher scale the same mid - uv spectral comparisons as are shown in figure 1 , plus that of the sun . panel @xmath33 shows 2627.5  2630.5 , dominated by strong lines . the other panels show regions of locally low line blanketing : @xmath38 , 2642.5  2645.5 ; @xmath39 , 2900.0  2903.0 ; @xmath29 , 3082.5  3085.5 ; @xmath30 , 3100.0  3103.0 . in all plots in fig .  3 , the spectrum of 184499 is offset by 30% from that of 128620 = , which in turn is offset by 50% from that of the sun , for which 1.0 marks the height of the true continuum . the same normalization constant as in fig . 1 has been adopted for each stellar spectrum , as was a normalization for the sun 1.10 times that given in the @xcite atlas , as noted on the left . identifications are provided at the top for the strongest lines as calculated in the solar spectrum . manually added lines are flagged with a colon following the decimal digits of the wavelength in  ngstroms . next appears the species , the lower excitation potential in ev of atomic transitions or the band and line of molecular ones , the digits following the decimal of the residual flux in the unbroadened line of the solar spectrum calculation , and the log gf - value adopted .    in fig . 3@xmath30 of the 3100.0  3103.0  region , calculated and observed spectra agree reasonably well in all stars , including  and the sun . no `` missing '' lines are strong enough to merit explicit identification at the top of the figure . true continuum is reached in each spectrum , but fleetingly , near 3102.5 ; its placement is sensitive to the treatment of adjacent weak lines . an unblended 0 - 0 oh line appears at 3101.230 , next to a near- continuum high point ; both are modeled well in all spectra . the strong atomic lines are well reproduced , as are two moderately strong ch lines at 3100.120 and 3100.170 . 3@xmath29 offers a further picture of how well both continuum and oh features are modeled . another continuum window appears near 3084.7 , but again with a weak line superimposed in the solar and  spectra . the relatively unblended oh 0 - 0 band line at 3084.896 , used in several oxygen - abundance investigations , is well reproduced over the factor - of-100 metallicity range of the stars in whose spectra it is detected , including the sun itself . nearby at 3085.2  is a blend dominated by a stronger 0 - 0 oh line ; it too is reasonably well matched in all stars . recall that to achieve this agreement , the @xcite theoretical gf - values for oh were reduced by @xmath40dex ; otherwise these features and that of fig .  2@xmath30 , plus the overwhelming majority of other oh lines , would be too strong in the sun . that is , unless the solar continuum were raised . the solar continuum appears to be well defined , and the oh lines of all bands matched to @xmath20.1dex , throughout the 3000  3100  region . because there are so few continuum windows , however , we can not rule out a systematic underestimate of the continuum in  and the sun of perhaps 5% or 10% . roughly speaking , raising the  mid - uv continuum would allow it to be modeled at an [ fe / h ]  value more consistent with that of @xcite , while raising the solar continuum might eliminate the present need to reduce the theoretical oh gf - values by 0.15dex . as noted at the end of  6 , it would also bring the mid - uv normalization constants of the sun and  closer to their expected values . a definite answer could be obtained by deriving the stellar parameters of  from an optical spectrum in the same way as for the other stars , and from modeling an improved spectrum of the solar near - uv . random continuum variations approaching 5% occur for two , and possibly all three , of the stars that were noted in  3 as observed through a smaller aperture . for the continuum seems appropriately placed for all stars in fig .   3@xmath30 , but in fig . 3@xmath29 the continuum of the 106516 observed spectrum appears too low , and in fig . 3@xmath39 that of 128167 may also be too low . such variations might also be present in the 128620 =  spectrum as well , but remain undetected because of the difficulty of matching its multitude of lines . @xmath41 shows another low - opacity window at 2900 . here the number of missing lines added remains small , and agreement is nearly as good . oh lines are less important , being generally minor constituents of blends . however , windows of true continuum have disappeared from the solar and  spectra , and are fleeting at best in the echelle spectra of lower resolution that were acquired for mildly metal - poor stars . continuum placement depends not so much on oh features but rather on whether all missing lines have been spotted . clearly that is not the case at solar metallicity , where weak missing lines apparently overlap to the extent that they can not be distinguished as individuals . although the cumulative effect of their neglect on the total absorption is small , is it significant for the continuum placement , which consequently can not be determined as well as was the case above from fits such as this . it is placed here and for the remaining panels by adopting the same normalization factors required to match figs . 3@xmath29 and @xmath30 . by 2645 , in fig .  3@xmath38 , the fits have deteriorated considerably . this is especially true at and above solar metallicity , where the numbers of missing and identified lines noted at the top are comparable . true continuum still appears in the weakest - lined spectra near the top of the panel , but elsewhere it is gone , lost to a forest of weak lines . although the fit to the metal - poor stellar spectra is still reasonable , the high - resolution spectra of the  and the sun reveal that too few missing lines have been added , often too strong and in the wrong places . for the adjacent region from 2645 to 2700 , the sheer number of missing lines precluded their reliable identification at the low resolution of the low - metallicity stellar spectra . because many were simply omitted , the calculations overestimate the spectral flux in that region  mildly in moderately metal - poor stars but seriously at and above solar metallicity , according to fig .  1 .    from fig .  3@xmath33 , it is clear that these difficulties are reduced in regions dominated by strong atomic lines . there the fit is dictated more by the ability of the calculations to match the cores and the damping wings of strong lines , and not by modeling the weak or missing features . these are largely suppressed : although the region shown is blueward of that of panel @xmath38 , only one missing line is noted in it . the wings of the strong lines are reproduced well in the four top spectra but begin to be somewhat too broad near the core in 106516 and 201891 . the entire @xmath31  about the core is too deep in   and the sun . this could perhaps be remedied by continuing the temperature enhancement to shallower surface layers than those indicated in  4 . figure 4 shows fits to the optical spectra of the sun and all stars except . identifications are given for the strongest lines in the solar spectrum . since the optical spectra were not fluxed , the continuum was normalized to unity . the same models were used to calculate the synthetic spectra as in figures 1 and 2 . the wavelength regions in panels @xmath33 @xmath30 depict a variety of magnesium lines and those of other heavier elements , to illustrate the extent to which the spectral calculations match features of a wide range of excitation , ionization , and strength . figure 4@xmath42 shows , to illustrate the satisfactory choice of . the normalization of the  profile was aided by its position near the center of an echelle order , and was done by interpolating the normalization functions of the two adjacent orders , centered @xmath375  away . this was straightforward , for the  line is much narrower than this and the continuum is very well defined by a spline fit of low order ; moreover , the normalization functions for the adjacent orders were very similar . note that most absorption lines in this panel are telluric , not stellar , as evidenced by their shift in apparent wavelength from one star to the next . included in fig . 4@xmath42 above the best fit for 94028 is the spectrum for the cooler 94028 model whose mid - uv spectrum is shown in fig . this illustrates how the fit to  deteriorates when models of 100  lower  are adopted . the  profile is relatively insensitive to gravity in this case : calculations adopting a model that differed from the best - fit model only in having  0.2dex larger produced a profile that deviated only near 6560  and 6565 , by about the thickness of the heavy line . similar calculations showed no sensitivity of the  profile to the chromospheric elevation of temperature , nor to changes of 0.3dex in iron abundance or light - element ratio . however , @xcite models required a  200  hotter , thanks to their inclusion of convective overshoot . panels @xmath33  @xmath30 of fig .  4 show that all the optical magnesium features are well matched using the same magnesium abundance and the same model as for the best fit to the mid - uv and to . panel @xmath33 plots the region near the very high - excitation mgii  doublet at 4481.2 ; @xmath38 , the ground - state mgi  line at 4571.1 ; @xmath39 and @xmath29 , the  lines at 5167.3 , 5172.7 , and 5183.6 ; and @xmath30 , the high - excitation mgi  line at 5711.1 . these panels also demonstrate that the calculations match strengths and profiles of other atomic lines regardless of species , excitation , or strength . among the lines depicted are tii  at 5173.7 , tiii  at 5185.9 , ground - state fei  at 5166.3 , feii  at 4576.3 , and sii  at 5708.4 . these panels also contain missing features : not modeled at all in the calculations are the absorption lines seen near 4480.8 , 5170.7 , and 5181.3 . they are not telluric , since they appear at the same position in the stellar spectra , and are especially visible in the two stronger - lined stars . our results for the parameters of the eight metal - poor turnoff stars analyzed here are listed in columns 12  15 of table 1 . the  values found for these stars by works cited in  1 are included for reference . our  values are seen to fall in the middle of the range of previous results , preferring neither the hotter nor the cooler values . unfortunately , the number of stars is too small and the scatter too large for more definitive conclusions . the reliability of our  values is supported by a comparison of model colors to those observed . for each star including the sun , we interpolated the @xmath43 and  colors and @xmath5 magnitudes calculated by @xcite to the , , and [ fe / h ]  values of each model we used . we have made no reddening corrections , nor any corrections for the temperature enhancement invoked in the more metal - rich models as described at the end of  4 . the ninth , tenth , and eleventh columns of table 1 show the difference in , , and  between the observed and calculated colors for each star . the agreement is extremely good in all colors over the entire metallicity range . the mean difference is @xmath440.010 , @xmath440.018 , and 0.000mag in , , and  respectively . the standard deviation of an individual difference is 0.010 , 0.026 , and 0.015mag , marginally larger than typical observational errors alone . the color differences are small even for , for which both metallicity and  are influential ( we interpolated to a model [ fe / h ]  = + 0.05 ) , and for which we have no optical spectrum and so did not derive parameters . for the other stars ,  and are sensitive primarily to @xmath45 ; their deviations correspond to @xmath345  with no allowance for observational error . their observed colors are slightly bluer than those of models , but the size of the differences points to a scale error in   of less than 50 . we thus estimate an uncertainty of @xmath350  in each  determination , from this and the poorer quality evident in figs .  2 and 4@xmath42 of the fits to mid - uv fluxes and  profiles as  is increased by 100 , plus the deterioration of the fit shown in fig . 4@xmath33 to mgii . from plots like those of fig .  4 at other wavelengths , we estimate the uncertainty in [ fe / h ]  to be @xmath30.05 dex at fixed , increasing to @xmath30.1 dex when the  uncertainty is included . the mid - uv slopes and  profiles are only weakly sensitive to , so our results for  are susceptible ( as are those of other groups relying on ionization equilibria ) to errors in gf - value scales and to possible non - lte effects . we have not yet taken full advantage of the technique of matching damping wings of strong lines , incorporating improved treatments of hydrogen broadening such as those of @xcite and @xcite . it would be valuable to do so given the constraints this would offer on alternative measurements of the fundamental properties of these standard stars , for example parallaxes from hipparcos and other space missions . however , the following comparisons of observed and model fluxes indicate that the  values are good to @xmath30.1dex as they stand . we first compare the uv normalization constant with that expected . according to @xcite , the ratio of observed to emitted stellar flux is @xmath46 , where the angular diameter @xmath47 = 2@xmath48/@xmath49 , @xmath48 being the stellar radius and @xmath49 its distance . for each star we calculate @xmath50 in solar units as @xmath51 , assuming masses @xmath52 = 0.8  for for the metal - poor stars based on stellar isochrones @xcite and 1.15  for 128167 , and adopting 1.16  for  as derived from its orbit by @xcite . the distance in solar units is found from the parallax , and the theoretical flux normalization constant follows from @xmath53 times 4.2545 @xmath54 to convert to square steradians . the ratio of this theoretical constant to that actually used to scale the observations in figs .  1 and 3 is listed in column 7 of table 1 . agreement is good for all stars but one . for , the ratio is 0.46 , suggesting an observed flux that is too low . indeed , @xcite saw this in their 1995 mid - uv  hst observations with ghrs : relative to earlier ghrs and iue observations of , `` our fluxes are lower ( by factors of 5.6 and 4.8 for the mgii  and feii  spectra , respectively ) , presumably because of the star s being misplaced at the edge of the small aperture . '' from a visual comparison with fig .  1 of @xcite , the fluxes we are using for the @xcite stis  spectrum should be raised by a factor of 2.1 . this brings the  ratio to 0.96 , and results in an average of the nine stellar ratios of 1.06 @xmath3 0.05 , with an individual standard deviation of 0.15 . for the sun , the ratio is 0.91 given the scale factor of 1.10 noted above , which is within the stated normalization uncertainty but still low . these  and solar ratios would be closer to unity were the continuum levels for their spectra drawn higher by 5%  10% , a possibility raised in  5 . as matters stand , errors in continuous opacity are limited to @xmath1110% nonetheless .    the apparent @xmath5 magnitude may also be compared to that expected from the model . from the @xmath5 magnitudes tabulated with the model colors , we calculated the difference in the stellar model @xmath5 with respect to the solar model , added this and @xmath55 to the observed solar @xmath5 = @xmath56 , then subtracted this from the observed stellar @xmath5 magnitude . again , no reddening or other corrections were applied . the resulting differences are shown in column 8 of table 1 . the agreement is reasonable : the theoretical magnitudes average 0.09 @xmath3 0.06mag brighter than those observed , with a @xmath3 0.18mag individual standard deviation . in brief , we have calculated high - resolution mid - uv spectra from first principles , after modifying line parameters and assigning line identifications for missing lines based on changes seen among spectra of stars spanning a range of well - determined values of @xmath45 , log @xmath35 , and [ fe / h ] . the resulting calculations provide a good fit for turnoff stars of all metallicities redward of 2900 . previous discrepancies in fitting the solar pseudocontinuum ( noted in  1 ) are attributed not to missing continuous opacity , which is good to @xmath1110% near 3100 , but rather to missing or miscalculated line opacity .    to improve the accuracy of the mid - uv gf - values , damping constants , and line identifications derived here , we have determined temperatures , gravities , microturbulence , and abundances for eight stars by calculating fits to optical as well as uv spectra . for stars more metal - poor than one - thirtieth solar , the fit to the mid - uv spectrum was satisfactory using laboratory lines alone , and can by itself determine . for the more metal - rich stars , lines missing from the laboratory list compromise the goodness of fit of the mid - uv spectrum . a 100  increase in  is seen to be difficult to detect from mid - uv spectra alone , as long as [ fe / h ]  and  are increased by 0.1dex to compensate . this -[fe / h ]  degeneracy parallels the age - metallicity degeneracy in elliptical galaxies @xcite . as with galaxies , it is reduced here by considering the balmer lines , with the advantage that stellar models need not include balmer - line emission nor composite populations . here  was set for stars with [ fe / h ]  @xmath57 by matching the  line profile and the strengths of other high - excitation features .    for all the metal - poor stars , the excellent match between calculations and observations of both optical and mid - uv spectra , and between model and observed colors , indicates an uncertainty of @xmath350  in an individual   determination , and a similar uncertainty in the  scale . moreover , the stellar @xmath5 magnitudes generally agree to 0.1mag , and the mid - uv normalization constants to 10% , with those expected from the stellar angular diameter found from the observed parallax and the model , assuming reasonable masses .    in this process , it was found that these optical diagnostics and the slope of the mid - uv continuum could only be fit simultaneously when using models of @xcite in which convective overshoot has been turned off . our work thus agrees with that of @xcite , who found that their models better reproduce the balmer profiles in the sun than do those of @xcite . downloading the @xcite models from the kurucz web site at http://cfaku5.harvard.edu is recommended for all work where the results depend upon the temperature structure in deep continuum - forming regions . residual discrepancies in fitting the profiles of very strong atomic lines , and the strengths of strong versus weak solar oh lines in the 3000  3100   region , are reduced here by mimicking a chromosphere towards the surface of every model with [ fe / h ]  @xmath58 @xmath59 . as described at the end of  4 , the model temperature was raised in the region just above the depths where optical continuum and weak lines are formed . while this procedure was necessary to match strong - line cores and oh line strengths over a very broad range of metallicity with a single set of gf - values , it is definitely an oversimplification . energy and stability considerations remain to be examined . fits might be improved with a different choice of temperature enhancement . indeed , a two - stream model might be found necessary . in the meantime , line cores also will be subject to uncertainties due to the presence of chromospheric emission . @xcite demonstrate that mgii  emission is always present in solar - temperature stars , with a strength that increases with increasing metallicity . emission also increases with increasing stellar activity @xcite . since activity in turn is associated with rapid rotation , stronger mgii  emission appears commonly among young stars and close binaries . consequently , no single quiescent chromosphere appears capable of modeling the very core of the profile of a strong line in all turnoff stars . the reliability of our  values and the good match achieved for oh lines enable a second look at the oxygen abundances of metal - poor stars , which are still uncertain as noted in  1 . our values for  are intermediate , suggesting that  is not the sole cause of the oxygen discrepancies . we defer a more detailed discussion to a future paper by @xcite , who compare oxygen abundances derived from the oh lines with those from the permitted and forbidden lines of atomic oi  in these stars . the manual addition of lines to the model line list , to match otherwise missing features seen in the observed spectrum , is clearly necessary for even an approximate representation of the mid - uv spectrum of solar - temperature stars of one - tenth solar metallicity and higher . nonetheless the procedure has its important limitations . one is that as overall line strength increases , the proportion of missing lines goes up dramatically ; disentangling them requires a broader range of spectra of high quality , as discussed below . another is that the excitation or ionization of the transition could be misjudged ; the ramifications are discussed in the next paragraph . a third is that many missing lines may not be due to iron as assumed , but to other elements . as discussed in  4 , this is potentially serious redward of the bound - free absorption edge of mgi  near 2512 . elsewhere the effect should be minimal , because although other iron - peak elements contribute significant line absorption throughout the mid - uv , the relative abundances of all other iron - peak elements with respect to iron change little from one late f or early g main - sequence star to another in both halo and disk populations in our galaxy . to be sure , the list generated by this procedure can not be expected to reproduce the spectrum of a peculiar a or f star , for example a cr - enhanced star , in which the iron - peak ratio has been dramatically altered @xcite .    more generally , the list is not suitable for modeling stars whose parameters lie well outside the range of those analyzed here . in such cases , any difference becomes important between the assigned and the actual excitation ( and ionization ) potential of the transition responsible for a line , leading to a significant misrepresentation of its strength . consequently , the present list can not be expected to reproduce the spectra of stars hotter than early f nor cooler than early g , nor giants , as well as stars of metallicity substantially higher than solar for which lines are still missing . the list could be made suitable for those stellar types by extending the present analysis to include such stars , provided high - resolution , high s / n mid - uv spectra are available in which mismatches induced by errors in the list can be visually identified and corrected . we discuss this further below . conversely , the list as it stands should do as well as indicated by fig .  1 and fig . 3 in reproducing the mid - uv spectra of any other late f or early g star of solar metallicity or less , provided a suitable model atmosphere is chosen . this holds because once the gf - value and damping constant of a line transition are set , elemental abundance and the boltzmann and saha equilibria at each point in the atmosphere fully determine line strength in quiescent stars such as these . dynamical effects can be ignored , and non - equilibrium effects ( if any ) are slowly - varying functions of spectral type , hence absorbed into the gf - value determinations . the only exception is the chromospheric emission noted above . were the line list improved to the point that a reliable match is achieved to the mid - uv spectra of f and g stars of solar metallicity and above , we could expect to use the mid - uv slope to at least constrain age and metallicity , regardless of its overall level or the appropriate light - element abundance ratio , as these could be incorporated in a grid of calculations covering the entire relevant parameter space . spectral indexes offer an alternative means by which mid - uv spectral calculations alone might constrain metallicity and  of a star or age of a stellar system . the principle is to construct indexes based on spectral regions dominated by features whose behavior breaks the age / metallicity degeneracy and elucidates the alpha enhancement , examples of which are the balmer lines and ratios of fei / feii  and mgi / mgii  line strengths . spectral indexes have indeed been used as diagnostics of age and metallicity in ellipticals by several of the groups mentioned in  1 . for example , @xcite have compared mid - uv and optical line indexes they calculated from kurucz flux distributions against mean values observed for the stars considered by @xcite . unfortunately , poor agreement was seen for all but two mid - uv indexes , the slope of the pseudocontinuum between 2600  and 3100 , and blended absorption near 2538 . we should be able to improve upon some of those index calculations immediately . as figs . 1 and 3 illustrate , the current spectral calculations should be reasonably reliable at all metallicities in spectral regions redward of 2900 , and for those dominated by strong lines blueward of that . this includes several indexes already in use , notably the mgii  and mgi  lines at 2800  and 2852  and line blends in the 3000 +  region . we plan to begin such index calculations shortly , and make them generally available . blueward of 2900 , however , our calculations are not yet reliable enough . 1 shows that our calculated spectra of solar - metallicity stars substantially overestimate the true fluxes in those wavelength regions not dominated by strong lines . somewhat ironically , our calculations thus suggest that neither of well - behaved @xcite indexes is currently reliable at near - solar metallicity . the problem undoubtedly is lines still missing from the line list . as seen in fig .  3 , missing lines often overlap in the solar spectrum to such an extent that they could not be individually discerned , and so their species and wavelength assignments became arbitrary . this was true especially in the 2550  2600   and the 2640  2710  regions , where we did not complete the process . this is the reason for the excess flux seen in the calculations at solar metallicity in these regions . this could be remedied were spectra obtained of the same quality as that of  for stars of the line strength of 184499 that span a range of properties , notably @xmath60 . for example , the spectra of 184499 and a hyades f star with  @xmath26500  should have about the same overall line strength : the hyades star s higher abundance , [ fe / h ]  @xmath2 + 0.15 , is offset by its higher @xmath45 at the rate of 0.1dex per 100 , as fig .  2 illustrates . because line crowding is much lower than in , comparing high - resolution mid - uv spectra of 184499 and a sharp - lined hyades f star could show where the strongest lines are that are missing at solar metallicity and suggest their identities from the changes observed with @xmath61 . the  and solar spectra could then be used to fill in the rest . analyzing an additional star closer in  to 184499 but with  = 0 would show explicitly which transitions are due to magnesium in the 2550   region , and whether oh has been properly characterized throughout the 2900  3100  domain .    at that point , the full arsenal of indexes representing the strengths of mid - uv features could be calculated from the spectra , and their behavior with , , [ fe / h ] , and  assessed explicitly . this is a major advantage of calculations from first principles over empirical libraries , which are confined to the spectra of stars available nearby . it is a major goal of this project . indexes should result with significant diagnostic power for the integrated - light spectrum of an old stellar system . we are indebted to s. allen of u.  c.  santa cruz for writing the initial scripts to convert the vax version of kurucz codes into unix versions . we are grateful to m. bessell of mount stromlo and siding spring observatories for communicating the @xcite and lifbase gf - value calculations and programs for their comparison against the kurucz values , to j. valenti of the space telescope science institute for providing the 128620 spectral reductions , to t. misch of lick observatory for obtaining the blue spectrum of 114762 , and to j. fulbright of the dominion astrophysical observatory for contributing his optical echelle spectra . we thank fulbright and r. p. kraft of lick observatory , r. l. kurucz of the harvard - smithsonian center for astrophysics , t. m. lanz and s. r. heap of the nasa / goddard space flight center , d. c. morton of the herzberg institute for astrophysics , and r. w. oconnell of the university of virginia for useful discussions . this work has made use of the simbad database , operated at cds , strasbourg , france . support for this work was provided by nasa astrophysics data program contract s-92512-z , stsci grants go-07395 , go-07402 , and ar-8371 , and nsf grants ast 99 - 00582 and ast-0098725 , to astrophysical advances , and by nasa adp grant nag 5 - 7104 , nasa ltsa grant nag 5 - 6403 , and stsci grant go-06607 to the university of virginia . linsky , j. l. , pagano , i. , valenti , j. , & gagne , m. 2000 , in _ cool stars , stellar systems and the sun : eleventh cambridge workshop _ garca lpez , r. j. , rebolo , r. , & zapatero osorio , m. , a. s. p. conf . ser . , 223 , in press                                        peterson , r. c. , & schrijver , c. j. 2000 , _ cool stars , stellar systems and the sun : eleventh cambridge workshop _ , ed .  garca lpez , r. j. , rebolo , r. , & zapatero osorio , m. , a. s. p. conf . 223 , in press                                      plots are shown comparing observed ( heavy line ) and calculated ( light line ) spectra for the nine stars with fluxed mid - uv spectra included in this study . to the right , the name of each star in bold is accompanied by the model parameters used for the corresponding calculation . at the left is the mid - uv normalization constant for each comparison . the same as fig .  1 , but on a larger scale in 3  spectral regions . at the bottom is the sun : two scans of the center of the solar disk are superimposed . at the very top , the strongest lines in the solar spectrum calculation are identified as described in the text . the same mid - uv normalization constants as in fig . 1 are adopted for the top nine stars . @xmath33 ) 2627.5  2630.5 . @xmath38 ) 2642.5  2645.5 . @xmath39 ) 2900 - - 2903 . @xmath29 ) 3082.5  3085.5 . @xmath30 ) 3100  3103 . comparisons as in fig .  1 are shown for optical spectral regions containing magnesium or balmer lines in eight of the stars plus the sun . each observed spectrum ( heavy line ) has been normalized to unity as described in the text . line identifications as described for fig . 3 are shown for the strongest lines in the solar calculation . @xmath33 ) mgii . @xmath38 ) low - excitation mgi . @xmath39 ) , @xmath29 ) the  lines . @xmath30 ) high - excitation mgi . @xmath42 ) . in this panel only , the third plot from the top represents a non - optimal choice of model parameters . it illustrates the degradation of the fit to  when a  is chosen 100  cooler than that of the plot immediately below . rcccccccccccccccccccc 19445 & + 8.06 & 0.45 & @xmath62 & 0.349 & 25.85 & 1.37 & @xmath63 & @xmath64 & @xmath64 & @xmath65 & 6050 & 4.5 & @xmath66 & 1.0 & 6007 & 5842 & 6050 & 6066 &  & 5825 + 84937 & + 8.33 & 0.39 & @xmath67 & 0.303 & 12.44 & 1.21 & + 0.09 & @xmath64 & @xmath64 & @xmath68 & 6300 & 4.0 & @xmath69 & 1.5 & 6314 & 6220 &  & 6350 & 6202 & 6375 + 94028 & + 8.22 & 0.47 & @xmath70 & 0.343 & 19.23 & 1.07 & + 0.25 & + 0.00 & + 0.00 & @xmath65 & 6050 & 4.2 & @xmath71 & 1.3 & 6047 & 5898 &  & 6060 & 5998 & 5900 + 106516 & + 6.11 & 0.45 & @xmath72 & 0.317 & 44.34 & 0.99 & @xmath1 & @xmath64 & @xmath73 & @xmath74 & 6250 & 4.3 & @xmath75 & 1.0 & 6221 & 6118 & 6208 & 6267 & 6233 & 6200 + 114762 & + 7.31 & 0.53 & @xmath76 & 0.365 & 24.65 & 0.96 & + 0.30 & @xmath77 & + 0.01 & @xmath78 & 5850 & 4.0 & @xmath79 & 1.0 &  & 5790 &  & 5941 & 5904 & 5800 + 128167 & + 4.46 & 0.36 & @xmath80 & 0.254 & 64.66 & 0.92 & @xmath81 & + 0.01 & @xmath64 & @xmath82 & 6850 & 4.3 & @xmath83 & 2.0 &  &  & 6707 & 6734 & 6737 & + 128620 & + 0.00 & 0.67 & + 0.23 & 0.438 & 742.2 & 0.46 & @xmath76 & + 0.00 & + 0.01 & @xmath84 & 5800 & 4.3 & + 0.15 & 1.0 &  &  &  &  &  &  + 184499 & + 6.62 & 0.58 & + 0.00 & 0.390 & 31.29 & 1.11 & + 0.02 & @xmath77 & @xmath77 & @xmath85 & 5750 & 4.0 & @xmath86 & 1.2 & 5734 & 5726 & 5750 & 5750 & 5773 & 5700 + 201891 & + 7.37 & 0.51 & @xmath87 & 0.358 & 28.26 & 0.92 & + 0.40 & @xmath77 & @xmath88 & @xmath78 & 5900 & 4.1 & @xmath89 & 1.0 & 6014 & 5817 &  & 5974 & 5918 & 5825 + sun & @xmath56 & 0.65   & + 0.13 &  &  & 0.91 &  & @xmath64 & @xmath63 &  & 5775 & 4.44 & + 0.00 & 1.0 &  &  &  &  &  &  +
determining the properties of remote globular clusters and elliptical galaxies using evolutionary population synthesis requires a library of reliable model stellar fluxes . 
 empirical libraries are limited to spectra of stars in the solar neighborhood , with nearly solar abundances and abundance ratios . 
 we report here a first step towards providing a flux library that includes nonsolar abundances , based on calculations from first principles that are calibrated empirically . 
 because the mid - ultraviolet spectrum of an old stellar system is dominated by the contribution from its main - sequence turnoff stars , we have started by modeling these . 
 we have calculated mid - ultraviolet spectra for the sun and nine nearby , near - main - sequence stars spanning metallicities from less than 1/100 solar to greater than solar , encompassing a range of light - element abundance enhancements . 
 we first determined temperatures of eight of the stars by analyzing optical echelle spectra together with the mid - ultraviolet . 
 both could be matched at the same time only when models with no convective overshoot were adopted , and only when an approximate chromosphere was incorporated near the surface of relatively metal - rich models . 
 extensive modifications to mid - uv line parameters were also required , notably the manual assignment of approximate identifications for mid - uv lines missing from laboratory linelists . without recourse to additional missing opacity , these measures suffice to reproduce in detail almost the entire mid - uv spectrum of solar - temperature stars up to one - tenth solar metallicity , and the region from 2900  to 3100   throughout the entire metallicity range . 
 ramifications for abundance determinations in individual metal - poor stars and for age - metallicity determinations of old stellar systems are briefly discussed , with emphasis on the predictive power of the calculations .