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water vapor is an important molecule for the chemistry of interstellar and circumstellar clouds . the 6@xmath8 - 5@xmath9 masing transition of h@xmath3o at 22 ghz , which arises from levels around 700k , has been used since its detection by cheung et al . ( 1969 ) to trace high excitation gas around star forming regions and evolved stars . the size of the emitting regions at that frequency is typically of the order of a few milliarcseconds ( a few 10@xmath10 cm ) . hence , no information has been obtained from this line on the role of h@xmath3o at large spatial scales . other h@xmath3o lines have been detected from ground or airborne based telescopes like the 3@xmath11 - 2@xmath12 transition at 183 ghz ( waters et al . , 1980 ; cernicharo et al . 1990 , 1994 , 1996 ; gonzlez - alfonso et al . 1994 , 1998 ) , the 4@xmath13 - 3@xmath14 transition at 380 ghz ( phillips , kwan and huggins 1980 ) , the 10@xmath15 - 9@xmath16 transition at 321 ghz ( menten , melnick and phillips 1990a ) and the 5@xmath17 - 4@xmath18 transition at 325 ghz ( menten et al . also , the 1@xmath19 - 2@xmath20 transition of h@xmath21o at 547 ghz has been observed by zmuidzinas et al . ( 1994 ) . among these lines only the 3@xmath11 - 2@xmath12 transition at 183 ghz has been used to map the emission of h@xmath3o at very large spatial scale ( cernicharo et al . 1994 , hereafter referred to as cer94 ) . the map of the orion molecular cloud shown in cer94 is 6 orders of magnitude larger than the size of the spots detected at 22 ghz and for the first time an h@xmath3o abundance estimate was derived for the different large scale components of the orion molecular cloud . the iso satellite has provided the opportunity to observe thermal lines of water in the middle and far - infrared ( see the reviews by van dishoeck 1997 ; cernicharo 1997 ; and cernicharo et al 1997a,1998 ) . mapping of the sgrb2 molecular cloud by cernicharo et al . ( 1997b ) has definitely shown that water vapor is an ubiquitous molecule in molecular clouds with an abundance of 10@xmath7 . maps of the emission of several h@xmath3o lines in orion irc2 have been obtained by cernicharo et al . ( 1997a , 1998 , 1999 ) . observations of the central position have been also obtained by van dishoeck et al . ( 1998 ) , gonzlez - alfonso et al . ( 1998 ) and harwit et al . however , the iso observations of h@xmath3o have drawbacks . in addition to the limited spectral resolution and the high opacity of the thermal lines of h@xmath3o , the poor angular resolution provided by iso in the far - infrared prevents any detailed study of the spatial structure and physical conditions of the h@xmath3o emitting regions .    an important aid in deriving h@xmath3o abundances could come from the observation of another masing transition of h@xmath3o with similar properties to those of the 183 ghz line . the 325 ghz line of h@xmath3o was observed by menten et al ( 1990b ) in the direction of orion irc2 and other molecular clouds . however , no maps were obtained . here we report the detection of extended water emission at 325 ghz and show that the h@xmath3o abundance is @xmath4 10@xmath5 in the * plateau*. the present data show the importance of ground - based observations of h@xmath3o in deriving the abundance of h@xmath3o in molecular clouds and in providing useful contraints on the physical conditions of the emitting regions . our ground - based observations provide much finer spatial resolution than iso or swas , and an estimate of x(h@xmath3o ) as accurate as that obtained from the extremely optically thick h@xmath3o lines observed in the submillimeter and far infrared domains . the observations were performed with the 10.4 m telescope of the caltech submillimeter observatory at the summit of mauna kea ( hawaii ) on april 1@xmath22 1998 . the receiver , a helium - cooled sis mixer operating in double - sideband mode ( dsb ) , was tuned at the frequency of the 5@xmath04@xmath1 line of h@xmath23o ( 325.152919 ghz ) . the h@xmath3o line was placed in the upper sideband ( usb ) to minimize atmospheric noise from the image sideband ( which was at 322.35 ghz ) . lines in the signal sideband are severely attenuated relative to those in the image sideband , due to the atmospheric h@xmath24 line . therefore it was necessary to check the sideband origin of a given line . the tuning frequency was shifted by 100 mhz to confirm that the central feature in the spectrum of orion - irc2 was the water vapor line . the backend consisted of a 1024 channel acousto - optic spectrometer covering a bandwidth of 500 mhz ( @xmath25v=1.1 kms@xmath26 ) . figure 1 shows the observed spectrum which matches very well the previous observation by menten et al . ( 1990b ) except for the line intensity ( see below ) . the pointing was determined by observing the same line towards the o - rich evolved star vy cma and was found to be accurate to 5  . the weather conditions were very stable during the observations with an atmospheric pressure and temperature of 620 mb and -1.4@xmath27c respectively . the relative humidity was measured to be 4 - 5% . the measured opacity from tipping scans at 225 ghz was @xmath40.025 . during the same night we performed broadband fourier transform spectroscopy ( fts ) measurements of the atmospheric absorption with the fts described in serabyn and weisstein ( 1995 ) . model calculations using the multi - layer atmospheric radiative transfer model atm ( cernicharo 1985,1988 ; pardo 1996 ) yielded an estimated precipitable water vapor column above the telescope of @xmath28200  @xmath29 m , which corresponds to a zenith transmission at 325.15 ghz of @xmath28 60% ( the corresponding value for the image side bande was @xmath28 87% , hence the line intensities for the image sideband features are overstimated by a factor @xmath28 2 ) . the heterodyne h@xmath3o data were calibrated using an absorber at ambient temperature . the calculated system noise temperature , for the signal sideband , was @xmath42100 k. the orion spectrum shown in figure 1 shows that the lines from the image sideband are weaker than the 325 ghz h@xmath3o line , i.e. , just the opposite of that occurring in the spectrum of menten et al . ( 1990b ) . rather than a variation of the maser emission ( the h@xmath3o line profile in figure 1 is identical to that shown in menten et al . 1990b ) we think that this difference is due to much better atmospheric transmission during our observations . we estimate that our intensity scale is correct to within 20 - 30% . the orion - irc2 map was carried out in position switching mode by using the on - the - fly procedure with an off position 5 away in azimuth . the spatial distribution of the 325 ghz emission is shown in figure 2 together with that of ch@xmath2oh ( from the image side band ) and the 183 ghz emission from cer94 . integrated intensity maps for selected velocity intervals are also shown in figure 2 .    in order to compare the line profiles of the h@xmath3o lines at 183 and 325 ghz we reobserved a few positions at 183 ghz with the 30-m iram telescope in january 1999 . the weather conditions were also excellent with a zenith opacity at this frequency of @xmath4 1 . the spectra , together with those obtained in 1994 , are shown in figure 3 . the observed 5@xmath17 - 4@xmath18 line profile towards the center position looks similar to that of the 3@xmath11 - 2@xmath12 line observed by cer94 ( see spectra in figure 3 ) . however , the antena temperature of the line is 20 times weaker and the line profile , although covering the same velocity range , is shifted towards the red . taking into account the different beam sizes of the iram-30 m telescope at 183 ghz and the cso at 325 ghz , and the extension of the emission in the latter line , we estimate that the main beam brightness temperature ratio , @xmath30=t@xmath31(183)/t@xmath31(325 ) is 10 - 20 if both lines were observed with a telescope of @xmath32 beam size . on the other hand , t@xmath31(183)@xmath33 k. both @xmath30 and t@xmath31(183 ) are well determined and can not be related to a calibration problem as the atmospheric conditions were extremely good during our observations at both frequencies . similar values for * r * , i.e. , @xmath341020 , are also found at other positions in the cloud ( see figure 2 ) and represent a real difference in the brightness temperatures of the two lines . only at position @xmath35=-12  , @xmath36=48  the peak temperature of the 5@xmath17 - 4@xmath18 transition approaches that of the 3@xmath11 - 2@xmath12 ( the 5@xmath17 - 4@xmath18 line is , however , narrower ) . the 325 ghz emission at this position presents a local maximum clearly visible in the velocity maps ( see figure 2 ) .    like the 3@xmath11 - 2@xmath12 line , the 5@xmath17 - 4@xmath18 transition is masing in nature . there are some narrow features at 325 ghz but with intensities of only a few k , i.e. , much weaker than those reported at 183 ghz by cer94 . these features agree in velocity with those found at 183 ghz . however , @xmath30 changes drastically from feature to feature , a fact that reveals the maser nature of the emission . outside the central region the lines are very narrow ( @xmath373 - 5 kms@xmath26 ) . some of the 3@xmath11 - 2@xmath12 narrow velocity components have antenna temperatures above 2000 k and are probably a few arcseconds in size ( cer94 , gonzlez - alfonso 1995 ) . the observations at this frequency performed in january 1999 clearly indicate a variation in the intensity of some of these features with respect to those of cer94 . however , at positions where the line is dominated by the plateau emission ( @xmath35=28  , @xmath36=-16  and @xmath35=12  , @xmath36=-16  for example ; see figure 3 ) and outside the central region ( @xmath35=-12  , @xmath36=48  ; figure 3 ) the line shape and intensity do not show any significant change between both epochs . a possible explanation for the 183 ghz and 325 ghz emission being spatially extended could be that it arises from many masing point like sources strongly diluted in the beam . this is ruled out by the results of cer94 where even the strong features at 183 ghz ( t@xmath382000 - 4000 k ) show indication of some spatial extent ( see above ) . the densities needed to reproduce the observed brightness temperatures of the maser spots at 22 ghz would result on a thermal or suprathermal 183 ghz line . consequently , if the 183 ghz emission was arising from the same region than that of the 22 ghz line very large column densities will be needed to reproduce the observed 183 ghz and 325 ghz intensities . in addition , the weak and extended emission observed at 183 ghz by cer94 clearly indicates the presence of water vapour coexistent with the molecular gas in the orion molecular ridge .    in order to understand the behavior of the two masing lines we have modeled the radiative transfer of the rotational levels of p - h@xmath3o for the physical conditions of the orion molecular cloud . the radiative transfer method is described in gonzlez - alfonso & cernicharo ( 1997 ) , and the model consists of a molecular shell with diameter of @xmath39 cm ( size of 15@xmath40 at 450 pc ) which expands at a constant velocity of 25 km s@xmath26 . collisional rates between water vapor and helium were taken from _ green _ , maluendes & mclean ( 1993 ) . the helium abundance was assumed to be 0.1 , and the rates were corrected to take into account the collisions between h@xmath23o and h@xmath23 . we calculated the statistical equilibrium populations of the lowest 45 rotational levels of para - h@xmath3o for different temperatures ( t@xmath41=100 , 150 , 200 and 300 k ) , column densities n(p - h@xmath3o ) , and volume densities ( n(h@xmath3)=@xmath42 , @xmath43 and @xmath44 @xmath45 ) . figure 4 shows the main beam brightness temperatures ( t@xmath31 ) ( as observed by a telescope of 15@xmath40 beam size ) for the 183 and 325 ghz para - water lines ( thin and broad lines , respectively ) , together with @xmath30 ( dashed lines ) . these t@xmath31 were computed from the integrated intensity by assuming that the spectral emission is gaussian - shaped . inspection of fig . 4 indicates that the line intensity ratio @xmath30 increases with n(p - h@xmath3o ) for low column densities ( which depend on t@xmath41 and n(h@xmath3 ) ) . both masers are unsaturated in these conditions , but the higher opacity of the 183 ghz transition makes this line more sensitive to variations of n(p - h@xmath3o ) . for higher values of n(p - h@xmath3o ) , the 183 ghz line becomes saturated and the exponential amplification of the 325 maser line yields a decrease of @xmath30 . finally , when both the 183 and 325 ghz lines are saturated , @xmath30 approaches a nearly constant value or even decreases below 1 for high n(h@xmath3 ) and low t@xmath41 . the maser at 183 ghz is quenched for these later conditions , although t@xmath31 can still remain above t@xmath41 due to the suprathermal excitation of the line ( see cer94 ) . even for relatively low column densities ( n(p - h@xmath3o)@xmath46 @xmath47 ) , low temperatures ( 100 k ) and volume density ( @xmath42 @xmath45 ) , the 183 ghz line has an intensity larger than 10 k ( see cer94 ) . however , the possibility of appreciable amplification for the 5@xmath17 - 4@xmath18 line is much more restricted than for the 183 ghz line , due to the high energy of the levels involved in the 325 ghz line ( @xmath48 k ) , and to the higher frequency and einstein coefficient of this transition . this fact explains the difference in spatial extent between both transitions , so that the the 325 ghz line is spatially restricted to the plateau while the 183 ghz line is in addition detected in the ridge ( cer94 ) .    the water vapor column density that fits the observed log * r*@xmath49 depends strongly on the assumed values of n(h@xmath3 ) and t@xmath41 . the higher n(h@xmath3 ) and t@xmath41 , the lower n(p - h@xmath3o ) that is needed to obtain an appreciable amplification of the 325 ghz line . 4 shows that * r*@xmath50 is obtained in different panels for n(p - h@xmath3o ) ranging from @xmath51 @xmath47 ( n(h@xmath3)=@xmath44 @xmath45 , t@xmath41=300 k ) to @xmath52 @xmath47 ( n(h@xmath3)=@xmath42 @xmath45 , t@xmath41=100 k ) . however , some of these models are not compatible with the observed intensities . for n(h@xmath3)@xmath53@xmath43 @xmath45 and t@xmath41@xmath54150 k , we find that * r*@xmath55 yields t@xmath31(325 ) in excess of @xmath56 k and t@xmath31(183)@xmath54@xmath57 k. both the observed intensities and * r * are only compatible with more moderate values of n(h@xmath3 ) and/or t@xmath41 . the physical reason is that , for high values of n(h@xmath3 ) and/or t@xmath41 , the collisional pumping of the 183 ghz line becomes so efficient that the emission in this line reaches high intensities for column densities that provide t@xmath31(325 ) of 50100 k. of course , the plateau may have regions with very high densities and temperatures ( which will give rise , for example , to the emission at 22 ghz and to the narrow spectral features observed at 183 and 325 ghz ) , but these will be much smaller than the observed size of the cloud . the widespread emission from the plateau we observe at 183 ghz and 325 ghz is only compatible with moderate values of n(h@xmath3 ) and t@xmath41 . in our models , the 22 ghz line will have intensities similar to those already observed in orion ( genzel et al , 1981 ) only for the highest kinetic temperatures and volume densities in figure 4 . the brightest spots at 22 ghz could be correlated with the narrow features at 183 ghz , and with the relatively weak lines at 325 ghz . brightness temperatures above 10@xmath58 k can be obtained for large column densities , t@xmath59@xmath54150 k and n(h@xmath3)@xmath5410@xmath58 @xmath45 .    lower limits for n(h@xmath3 ) and t@xmath41 can be obtained from the observations of other molecular lines ( e.g. , blake et al . 1987 ) , so that we adopt n(h@xmath3)@xmath53@xmath43 @xmath45 and t@xmath41=100150 k. for these conditions we obtain n(p - h@xmath3o ) in the range @xmath60@xmath61 @xmath47 , and hence n(h@xmath3o ) in the range @xmath62@xmath52 @xmath47 . the water vapor abundance can be derived from the co data taken with similar angular resolution ( see cer94 ) . for the intermediate velocity wind cer94 derived a co column density of 10@xmath63 @xmath47 . hence , the x(h@xmath3o)/x(co ) abundance ratio in the plateau is around 1 , i.e. , x(h@xmath3o)@xmath28@xmath64 . in the ridge molecular cloud the 325 ghz is very weak . our models and the 183 ghz data provide an estimate for x(h@xmath3o ) of a few 10@xmath6 - 10@xmath7 ( see cer94 ) which is in good agreement with our iso results ( see cernicharo et al . 1998 , 1999 ) . the comparison of several masing transitions arising in relatively low energy levels of h@xmath3o allows us to constrain the physical conditions of the different emitting regions . so far , ground - based observations of these transitions with large radio telescopes are the only means to obtain the spatial distribution of h@xmath3o in interstellar clouds . j. cernicharo and e. gonzlez - alfonso acknowledge spanish dges for this research under grants pb96 - 0883 and esp98 - 1351e . pardo gratefully acknowledges the financial support of the _ observatoire de paris - meudon _ , _ cnes _ and _ mto - france_. the cso is supported by nsf contract # ast-9615025 .
we present observations of the 5@xmath04@xmath1 transition of water vapor _ at 325.15 ghz _ taken with the cso telescope towards orion irc2 . 
 the emission is more extended than that of other molecular species such as ch@xmath2oh . 
 however , it is much less extended than the emission of water vapor at 183.31 ghz reported by cernicharo et al ( 1994 ) . 
 a comparison of the line intensities at 325.15 ghz and 183.31 ghz puts useful constraints on the density and temperature of the emitting regions and allows an estimate of h@xmath3o abundance , x(h@xmath3o ) , of @xmath410@xmath5 in the plateau and @xmath410@xmath6 - 10@xmath7 in the ridge . 
 21.75 cm -0.5 cm = 20pt      ism : molecules  ism : individual ( orion irc2 )  line : profiles  masers  radio lines : ism  submillimeter