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the different evolutive stages of massive stars have strong influence on their surroundings and , consequently , on the galactic evolution . during their lives , the massive stars inject large amounts of matter , energy and momentum into the interstellar medium ( ism ) . in the pre - main sequence , massive stars undergo a phase of very energetic mass loss giving rise to molecular outflows . uv radiation , expansion of hii  regions and fast stellar winds are present in the main sequence stage . in the short  lived red supergiant ( rsg ) and luminous blue variable ( lbv ) stages , the winds become slower and denser and ejections of up to some solar masses may be produced . the wolf  rayet ( wr ) stage , prior to sn explosions , is characterized by a copious mass loss ( typically @xmath5 m@xmath6 yr@xmath3 ) driven by a fast ( about 2000 km s@xmath3 ) and chemically enriched stellar wind . massive stars are therefore continuously heating , ionizing , shocking and blowing - up the ism surrounding them . garca - segura & mac low ( 1995 ) have modeled the evolution of the gas which surrounds massive stars , from the main sequence to the wr stage . they only considered simple models for the mass ejection in the form of stellar winds in the different stages of evolution and predict the presence of large scale structures around the massive stars . both the blown - up gas and the ejecta form multiple shells as a consequence of the different shockfronts produced inside the gas .    during the main sequence phase , the gas around massive stars is ionized , heated and highly evacuated by the fast , long - lived stellar wind and the lyman continuum radiation . a hot cavity is surrounded by an expanding shell . the ionization front may get trapped within it , the shell may start to recombine and become observable in the line radiation emitted by neutral atoms or molecules . these large bubbles have been successfully detected in the 21-cm line of atomic hydrogen ( @xcite ; @xcite ; @xcite ; @xcite ; @xcite ) , in the form of huge shells of 3070 pc , expanding at 620 km s@xmath3 , with masses of hundreds ( or even thousands ) of solar masses . some of the large iras shells found by marston ( 1996 ) might also be related to this type of structures . once the massive star finishes the hydrogen burning in its nucleus and becomes a rsg or lbv , its stellar wind suddenly becomes more dense , and a significant amount of the stellar mass is deposited into the ism at velocities around 20 km s@xmath3 . very probably , the formation of the optical ring nebulae starts in this brief stage . when the star comes to a wr stage , the wind is accelerated to hypersonic velocities and rapidly reaches the rsg / lbv wind . many of the optical nebulae possess arcs and have high inhomogeneities in the abundances of heavy elements ( @xcite , 1999 ) .    however , the effects of the wr phase in the neutral ( neither atomic nor molecular ) gas have not been easily detected . the structures formed during this phase would not be as large as those developed in the o - phase . just in few cases , the neutral atomic hydrogen has been detected in interaction with ring nebulae ( @xcite ; @xcite ) . vibrationally excited molecular hydrogen has only been detected towards ngc2359 . the co j = @xmath0  emission has been observed around wr16 ( @xcite ) and in ngc2359  ( @xcite ) . ngc2359  is a very interesting object , and many aspects of which has been thoroughly studied in the recent years . this optical nebula is excited by hd56925(wr7 in the catalogue of van der hutch et al . ( 1988 ) , a wn4 star located in the outer milky way . its distance varies from 3.5 to 6.9 kpc depending on the authors ( see @xcite for a discussion ) . the mass  loss rate ( @xmath7 m@xmath6 yr@xmath3 ) and the terminal wind velocity ( 1545 km s@xmath3 ) have been measured by schneps et al . ( 1981 ) and rochowicz & niedzielski ( 1995 ) , respectively . ngc2359 , a prototypical wind - blown bubble ( @xcite ) , is nearly spherical with several small filaments inside . the nebula is not highly enriched of heavy elements as a whole , with the exception of the filaments ( @xcite ) . this may indicate that the nebula was largely produced in the rsg / lbv and the wr stages of the exciting star . from the optical emission lines ( @xcite ) and radio recombination lines ( @xcite ; @xcite ) , we know that the ionized gas has several components at velocities between 30 and 70 km s@xmath3  ( all velocities in this paper are referred to lsr ) . an hii  region partially surrounds the wind - blown bubble and is apparently limited by obscuring material . schneps et al . ( 1981 ) have taken several profiles in the j = @xmath0line of co and detected the presence of three clouds at 37 , 54 and 67 km s@xmath3 . the presence of broad co profiles at 54 km s@xmath3  and the morphology of the emission , just following the south - eastern edge of the nebula , suggests that this component has been interacting with the nebula . st - louis et al . ( 1998 ) have analyzed the h@xmath8  emission in this region . they have found the 10 s(1 ) line towards the southern border of the hii  region , but they could not establish the nature of the excitation ( fluorescence or shocks ) of the h@xmath8 . cappa et al . ( 1999 ) made a complete map of the ionized component at the 1465 mhz continuum , and traced the hi  emission presumably connected with the nebula . these hi  features are probably associated with the wr stage of hd56925 ; so far , the large atomic main - sequence bubble predicted by the theory has not been yet reported . the goal of this paper is to look into the history of the interaction of the wind and the radiation of hd56925  and its surroundings . we have analyzed a public hi  survey and studied for the first time the large scale interaction of hd56925  in the o - phase with the ambient ism . based on the mapping of the j = @xmath0  and j = @xmath1  lines of co over the whole optical nebula , we study in detail the morphology and kinematics of the co around ngc2359 . we also present several @xmath4co profiles towards selected positions in order to determine the global physical parameters of the molecular gas , such as kinetic temperature , density and opacity of the lines . this study has allowed us to shed some light on the dominant excitation mechanism of the neutral gas in the wr stage . we have analyzed a region of @xmath9 around ngc2359  using the 21 cm  hi  survey of hartmann & burton ( 1997 ) . this survey was made using the 25 m radio telescope at dwingeloo . the angular and velocity resolution of this dataset are 36  and 1.03 km s@xmath3 , respectively . the survey has been done with a sampling of 30  in galactic coordinates and have a sensitivity in brigthness temperature of 70 mk .      the observations of the co j = @xmath0  ( 115.271 ghz ) and j = @xmath1  ( 230.538 ghz ) lines and the j = @xmath0  ( 110.201 ghz ) line of @xmath4co towards ngc2359  and its surroundings were made during 2000 february , using the nrao 12 m radio telescope at kitt peak . the half power beam width of the telescope was 54  and 27  at the rest frequency of the j = @xmath0  and the j = @xmath1lines , respectively . the data were acquired simultaneously with two independent sis receiver channels with orthogonal polarizations . the receiver temperatures were 250300k and 500600k for the j = @xmath0  and the j = @xmath1  lines , respectively . we used filterbanks ( with 256 channels ) and the hybrid correlator mac ( in the 2 if mode , in parallel , 3072 channels ) as backends . in the j = @xmath0line , the filterbanks provided 100 and 250 khz of resolution , corresponding to velocity resolutions of 0.26 and 0.65 km s@xmath3 , respectively . for the same line , the mac correlator was setup with a total usable bandwidth of 300 mhz and with channels of 195.3 khz of resolution ( 0.51 km s@xmath3 ) . in the j = @xmath1line , the filterbank with channel widths of 250 khz ( 0.33 km s@xmath3 ) and 500 khz ( 0.65 km s@xmath3 ) were used . we also used for this line the mac correlator tuned with an usable bandwidth of 600 mhz and frequency resolution of 390.6 khz ( 0.51 km s@xmath3 ) . mapping of the j = @xmath0  and j = @xmath1  lines was made using the `` on - the - fly '' technique . we have obtained three individual maps in the j = @xmath0  line and six in the @xmath1  line . these individual maps were later reduced with aips and combined weighting them by their inverse square of _ rms_. for our investigation , we finally used the mac data for the j = @xmath0  line and the 500 khz - filterbank data for the j = @xmath1  line . typical final _ rms _ per grid point was 120 mk and 80 mk in the j = @xmath0  and j = @xmath1  lines , respectively . we have also observed the j = @xmath0  line of @xmath4co and c@xmath10o at seven individual positions , with 710 minutes of integration time achieving a _ rms _ noise of 30 mk . the pointing was checked at the begining of each observing session with five  point cross maps . the final data are pointed within 4  of uncertainty . the data were calibrated using the chopper wheel method . final intensity scale is in units of @xmath11 , e.g. , corrected by atmospheric attenuation and instrumental losses ( @xcite ) . the calibration was also checked off - line by comparing individual spectra of standard sources . we have analyzed the hi  data in the velocity range 2080 km s@xmath3 , e.g. , the velocity range of the optical and the co emission . according to most of the galactic rotation models ( for example @xcite ) , this broad range in velocity is expected for the range of distances associated to hd56925(@xcite ) . within this velocity range , we looked for the structures connected to the evolution of hd56925 . 1 depicts hi  column density ( n@xmath12 ) maps as a function of the radial velocity , indicated at the top left corner . the white circle near the center of the map sketchs the position and size of ngc2359 . every map shown in fig . 1 has been obtained by the integration of the brightness temperature over 2.1 km s@xmath3  ( 2 channels ) and assuming optically thin emission . we have found an outstanding structure which remains coherent across tens of velocity channels . this structure is a huge cavity surrounded by a set of relative maxima in the n@xmath12  maps . the centroid of this cavity remains approximately at the same point , and nearly off - centered from the nebula and the wr star . the excentric position of the wr star with respect to the cavity is also observed in almost all the wr bubbles ( @xcite ) and might be due to the star motions and/or density variations in the original ism ( @xcite ) . the structure reaches its maximum size at velocities between 61 and 65 km s@xmath3 . for receding and approaching velocities with respect to this velocity interval , the cavity diminishes its angular size , ending at @xmath13 51 and 76 km s@xmath3  with small relative maxima centered at the cavity . the morphology of this structure has the typical ring disk appearence of an expanding bubble . following the standard analysis of this type of shells , the largest size of the cavity indicates the geometrical dimensions and its radial velocity corresponds to approximately the systemic velocity of the gas before the expansion . the full velocity range where the feature is observed gives a lower limit of twice the expansion velocity . table 1 summarizes the parameters derived , as described above , for the expanding cavity . following arnal et al . ( 1999 ) , we have also derived lower and upper limits to the mass of the hi  shell . these two values for the hi  mass are the missing mass into the higher close contour and the neccesary mass to fill the cavity making it indistinguishable of the nearly environs . all distance - dependent parameters were calculated assuming a distance of 5 kpc . this hi  shell is in geometrical coincidence with the iras shell reported by marston ( 1996 ) . furthermore , the mass of gas predicted from the iras data falls within the hi  estimates presented in table 1 .        the co emission in the j = @xmath0  and j = @xmath1 lines are observed in three velocity components . the components with the lowest radial velocity ( 3538 km s@xmath3 ) and with the highest radial velocity ( 6469 km s@xmath3 ) show narrow profiles of 1.52.5 km s@xmath3  width . hereafter we will refer to these two components as ambient 1 ( a1 ) and ambient 2 ( a2 ) , respectively . the third component , with radial velocities of 5058 km s@xmath3 , shows broader profiles than the components a1 and a2 , having linewidths of 45.5 km s@xmath3 . we will refer to this component as broad component ( b ) . 2 shows the global spatial distribution of the co j = @xmath0  ( left panels ) and j = @xmath1  ( right panels ) integrated emission for the three components . for comparison , the co maps have been overlaid with a j band image of ngc2359 , obtained from the _ digitized sky survey_. at the top right corner of each panel , the velocity interval of integration is indicated . the mapped regions in both spectral lines are also depicted . in the j = @xmath0line , almost all the optical nebula and its surroundings has been covered by our observations . in the j = @xmath1 line , however , we limited the maps at the east and the north - west because of our interest in studying in detail the interaction of the nebula with the molecular gas .    the co emission for components a1 and a2 appear mainly to the east and to the south - east of the mapped regions , respectively . although the spatial distribution for both components are found close in projection to the optical nebula , there is not a clear morphological correlation . furthermore , there are no significant changes in the line widths and in the peak velocities over the regions where these components are observed . in contrast , the broad component b is observed in clear correlation with the eastern and the southern edges of ngc2359 , and with broader profiles , varying from 4 to 5.5 km s@xmath3  wide . the most intense co emission of component b is located beyond the outer edge of the hii  region in ngc2359 , mainly to the south - east part , but also significant emission in the j = @xmath1  line is detected towards locations projected on the hii  region ( middle right panel in fig .  2 ) . for a comparison between the spatial distribution of the j = @xmath0  and @xmath1  lines of co for the component b , the fig . 3 shows and overlay of the emission in the j = @xmath0  line ( solid lines ) and the j = @xmath1  line ( dashed lines ) convolved to the resolution of the j = @xmath0  line . this comparison shows that both lines are well correlated at the most intense and the external part of the nebula . however , we note again that the j = @xmath1  line have significant extended emission to the inner part of ngc2359 , especially into the southern optical filament . this is clearly ilustrated in the left and right panels in fig .  3 , where we show a sample of four co j = @xmath0  and @xmath1  line profiles toward the positions sketched by squares in the central map . the lower left spectra show the mean profiles around the most intense region , while the other three panel show the co profiles in the periphery of the co emission . there is a clear trend in the profiles . while the j = @xmath0  line is stronger than the j = @xmath1  line toward the most intense zone of emission , the j = @xmath0  line is much weaker than the j = @xmath1  line or it is not detected in positions closer to the hii  region . the @xmath1/@xmath0  intensity ratio varies from 0.81 near the peak to 1.52 towards the positions near the hiiregion . in contrast , the @xmath1/@xmath0  intensity ratio is nearly constant in components a1 and a2 , with values around 0.50.6 . while the components a1 and a2 do not have any remarkable kinematical features , the component b have slightly different peak velocities as a function of the position . fig . 4 shows four position - velocity diagrams , taken in the directions sketched in the map of the j = @xmath1emission at the top left corner of this figure . each position - velocity diagram is identified by the slice number . slice 1 shows the most intense part of component b at 5354 km s@xmath3  and the western border of the component a2 . slice 1 also shows hints of a small velocity gradient in the component b. this is more clearly depicted in the slice 2 , where we find along the central part a velocity shift of @xmath13 2 km s@xmath3 . both slices 1 and 2 remarks the spatial coincidence between the eastern border of component b and the western border of component a2 . in fact , slice 3 shows the presence of a weak ( 4 @xmath14 ) `` bridge '' in the co emission which connects the components b and a2 . this suggests a physical connection between both components . slice 3 , parallel to the slice 2 but 1  to the south - west , does not show signs of this `` bridge '' , indicating the small size and weakness of this bridge . the slice 4 show that the diffuse emission to the north have a radial velocity slightly lower than that of the main part of the component b.      in order to study the excitation of the co line in the component b , we have observed seven positions in the @xmath4co and c@xmath10o j = @xmath0  lines , towards the southern part of ngc2359 . the location of the observed points are shown in the fig . 5a , superimpossed on a co j = @xmath1  map of the component b. five of these positions are located close to the slice 1 . 5b , 5c and 5d plot three of these @xmath4co j = @xmath0  profiles ( labeled with letters b , c and d in the circles of fig . 5a ) , together with the @xmath15co j = @xmath0spectra of the same points . we find a significant variation in the @xmath15co/@xmath4co ratio . this ratio is @xmath13 6 for the co emission peak ( fig . 5b ) and increases to values @xmath16 elsewhere . another interesting result is that the @xmath15co/@xmath4co ratio does not only vary with position but also it changes with radial velocity within the beam towards the co emission peak ( fig . one can see that both lines peak at different velocities . the @xmath4co line appears shifted by @xmath13 12 km s@xmath3  with respect to that of @xmath15co . this shift produces important changes in the @xmath15co/@xmath4co ratio as a function of the radial velocity . the @xmath15co/@xmath4co ratio is @xmath17 30 at 5053 km s@xmath3 , and decreases to @xmath13 45 for radial velocities between 54 and 57 km s@xmath3 . in view of the small velocity gradient found in the component b , we do not think that this shift can be due to a pointing error between the @xmath15co and @xmath4co data . even in the worst possible case , 10  of pointing difference , we find that this ratio would change with radial velocity by at least a factor of 4 for velocities below and above 53 km s@xmath3 . the hi  structure found in this paper resembles the expanding shells often associated with o stars ( @xcite ) , wr stars ( @xcite ; @xcite ) and ob associations ( @xcite ; @xcite ) . in the following we shall explore if the hi  shell has been created by the stellar wind from hd56925  or from its o  progenitor . let us estimate , with a few assumptions , some important parameters to distinguish between the two possibilities . we shall consider a spherical expanding bubble , with a radius of 50 pc , an expansion velocity of 12 km s@xmath3 , and a mass of 1500 m@xmath6 . with these values , we estimate a dynamical time for the shell of 2.3 myr and a kinetic energy of @xmath18 erg . the dynamical time was computed as a mean value between the adiabatic and the radiative case ( @xcite ) . this time is in fact an upper limit , but a good approximation , to the age of the hi  structure ( @xcite ) . this age and the kinetic energy associated to the hi  shell clearly indicate that its origin is due to the main - sequence phase of hd56925 . indeed , the most massive stars , greater than 25 m@xmath6 , remains a few million years at the main sequence ( @xcite ) and creates interstellar bubbles of 3060 pc in this stage ( @xcite ; @xcite ) . it is thought that wr s are descendant of these massive stars . on the other hand , the rsg , lbv or wr phases are short - lived , probably @xmath19 yr ( @xcite ; @xcite ) . we therefore conclude that only the o progenitor of hd56925  had sufficient time to develop the hi  expanding shell reported in this paper . a similar conclusion was reached by marston ( 1996 ) concerning to the large iras shell found around several wr stars , including the one linked to hd56925 .    by adopting typical values for the mass - loss rate of @xmath20 m@xmath6 yr@xmath3  and a wind velocity of 1000 km s@xmath3  for the o progenitor star , we estimate that this star deposited into the ism during 2.3 myr a total energy of @xmath21 erg in the form of stellar wind . if this star has blown up the shell which we are observed in hi , it implies a kinematical efficiency of nearly 10% . although the uncertainties in these computations are large and hard to be estimated , this value for the efficiency is in good agreement with the models that predict this type of structures ( @xcite ; @xcite ) and reinforces our hypothesis of a main sequence origin of the hi  shell . we have estimated the physical parameters of the molecular gas from the co line emission by modelling the excitation of the co lines using the _ large velocity gradient _ ( lvg ) approximation . for comparison with the model , we have smoothed the j = @xmath1  line data to the j = @xmath0  resolution . for component b , the @xmath4co data was also taken into account in order to estimate the opacity of the @xmath15co lines . six different components have been defined . the first three are the ambient components a1 and a2 and the most extended emission of the component b ( hereafter b@xmath22 ) . the fourth component considered is the one projected onto the southern hii  region , named as b@xmath23 . finally , we also considered the region around the co maxima in component b , where the @xmath4co behaviour is striking . we have subdivided this component in the two velocity ranges having the greatest difference in the @xmath15co/@xmath4co ratio ( see fig . 5b ) , named as b@xmath24 ( for the velocity interval 5054 km s@xmath3 ) and b@xmath25 ( for the velocity interval 5458 km s@xmath3 ) . the results derived from the lvg analysis are shown in table 2 . this table gives estimates for the excitation temperature ( t@xmath26 ) , h@xmath8  density ( n(h@xmath8 ) ) , co column density ( n(co ) ) , the opacities of the co j = @xmath0 , co j = @xmath1  and @xmath4co j = @xmath0  lines , and the total mass of the regions ( m(h@xmath8 ) ) , derived from the two line ratios ( second and third columns ) and the @xmath4co intensities . a kinetic temperature ( t@xmath27 ) of 10k was assumed for all the regions but b@xmath23 . the phyisical conditions in b@xmath23 seem to be very different from those of the other regions . the large ratio between the j = @xmath1  and the j = @xmath0  lines is explained by optically thin emission with relatively high t@xmath27 . the most likely t@xmath27  is about 80k and we present in table 2 the results derived for this t@xmath27 . components a1 and a2 are subtermally excited ( t@xmath26  @xmath28 t@xmath27 ) , with typical t@xmath26  of 56k . we obtain higher excitation temperatures in component b. for b@xmath24 and b@xmath25 , we derived t@xmath26  similar to t@xmath27 , indicating that the lines are thermalized , and densities significantly higher than for a1 and a2 . for b@xmath25 we did not apply the lvg method since this region has a very low value of the @xmath15co/@xmath4co ratio . such low value for this ratio usually indicates the presence of a region of very high opactiy in the @xmath15co lines . since we have found this low ratio toward only one observed position , we conclude that the emiting region corresponding to b@xmath25 is unresolved by the j = @xmath0  beam . to derive the physical properties in this component , we have assumed an optically thin emission for the @xmath4co and optically thick emission of the @xmath15co . for a @xmath15co/@xmath4co isotopic ratio of 60 ( @xcite ) , we derive an optical depth of 12 for the j = @xmath0  line of @xmath15co . from the @xmath15co j = @xmath0 line intensity and assuming that the emission in this line is thermalized to the t@xmath27  of 10k , we derive a beam filling factor for b@xmath25 between 0.1 and 0.2 . with the optical depth we have estimated n(co ) . adopting a beam filling factor of 0.15 and a fractional abundance for co of @xmath29 we have also estimated n(h@xmath8 ) . although our estimates of the h@xmath8  density can be uncertain because of the unknown t@xmath27 , the large differences in the h@xmath8  density between b@xmath25 and the other components indicate that a rather small region has a density larger than other components .      in sect . 3.2 we have discussed the morphological and kinematical arguments for the relationship between the different co components and ngc2359 . in the previous section , we have derived their physical conditions . in this section , we shall consider our observational findings together with those already published , in order to address the nature and the origin of the interaction of the molecular gas with the nebula and its exciting wr star . the component a1 is located mainly to the ne of the map , while the component a2 is present to the se of the nebula . both components are narrow ( 2 km s@xmath3 ) and do not show any morphological or kinematical structures that indicate any disturbance produced by the nebula or the wr star . in contrast , the component b has some kind of ubiquity in this field , in agreement with optical emission lines like the [ ] . it appears as a broad component ( 4.55 km s@xmath3 ) following the east and the south border of the hii  region which surrounds the nearly - spherical bubble . we noted that the co emission is stronger _ outside _ the optical border of the hii  region , but significant emission is also present in projection _ onto _ the hii  region . we think that the component b is the only component which undoubtely interacts with ngc2359 , because of its morphology , the large width of the lines , the higher temperature and the striking kinematics . as already mentioned , the eastern part and the southern border of the hii  region are clearly bounded by the most intense component b. furthermore , it is also striking the presence of faint co emission perfectly correlated with the hii  region . these results are also in agreement with the vibrationally excited h@xmath8  emission observed at 2.2@xmath30 m ( @xcite ) . the co emission is dense and more opaque outside the hii  region , while onto the hii  region it is less dense and has a higher temperature ( up to 80k ) . this temperature is compatible with the dust temperature inferred by marston ( 1991 ) in the same region .    at this point we wish to consider the @xmath4co behaviour in the southern part of the component b. we have only detected this molecule near the peak of the co emission . this indicates that this is a _ unique _ region where the @xmath15co/@xmath4co changes appreciably , most likely due to an increase of the opacity in the @xmath15co lines . only 1  to the east of this position ( essentially one beam ) there is no significant emission in @xmath15co from the component b , just at the western border of the component a2 . this spatial anti - correlation between components a2 and b is also present toward the southern interface . furthermore , there is a bridge in the co emission which seems to connect both components within a small area ( see slice 3 of fig .  4 ) , suggesting that both components are related . the co in the hii  region might be heated by the wr radiation field ( @xcite ) , but the interface between the components a2 and b may not be excited by the same mechanism . all these observational facts can be explained in a scenario in which a shock of 1014 km s@xmath3 , driven by the expanding hot bubble , impacts on the component a2 . as a consequence of the shock , the gas was accelerated up to the velocities of the component b. the large width of the component b when compared with the component a2 can be understood by means of an increase of the turbulent motion due to the energy injected by the shock . furthermore , the systemic velocity of the hi  shell found at larger scales ( presumably near the rest velocity of the gas ) is also similar to the velocity of the component a2 , supporting this scenario . the shock heats the gas and forms a _ thin layer _ of compressed material with large h@xmath8  density and with high opacity in the @xmath15co . this thin shocked layer would remain unresolved by the beam . the kinematics shown in fig .  4 and the low @xmath15co/@xmath4co ratio at 5356 km s@xmath3  towards the co peak are very suggestive of this thin shocked layer . indeed , the large optical depth of the co lines can only be explained if this emission arises from a very small region unresolved by the beam , probably less than 0.2 pc wide . furthermore , the large co column density of b@xmath25 as compared with b@xmath24 would indicate large densities for these radial velocities ( up to several 10@xmath31 @xmath32 ) . obviously , in the scenario of the shocked gas responsible for the component b , b@xmath25 would represent the material which has been recently shocked . high angular resolution observations are needed to disclose this particular region and to check our interpretation . although our hypothesis seems reasonable , it should be pointed out that this shocked region is small and it can not be spatially resolved in our data . another crucial point is the weakness of the `` bridge '' which connects the components a2 and b. higher angular resolution and more sensitive observations are needed to confirm the presence of the shocked layer and to fully study its properties . if this is the case , ngc2359  would have the first direct evidence of shocked molecular gas in a wr environment . the o - stage of hd56925  have created in the ism the hi  shell reported in this paper and the iras shell reported by marston ( 1996 ) . in consequence , the optical nebula and its surroundings , where we find most of the co emission , have a more recent origin than the large scale hi  shell . the hii  region has a size that indicate a dynamical age of less than 10@xmath33 yr . this age is comparable with the wr or a previous phase , such as lbv or rsg ( @xcite ) . on the other hand , many optical shells of enriched material are found close to the wr star ( @xcite ) . this leads us to think that the material now interacting with the hii  region was ejected in previous episodes of rsg or lbv phase . chemical studies using other molecular line transitions , especially those capable of tracing the molecular gas enriched by these phases , will definitively confirm or discard these ideas . we have analyzed the large - scale 21 cm - hi  emission in a @xmath34 field around the nebula ngc2359 , and found a large expanding bubble roughly centered at the nebula and its exciting star , the wolf - rayet hd56925 . the systemic velocity of the bubble is @xmath13 64 km s@xmath3  and it is expanding at 12 km s@xmath3 . at an assumed distance of 5 kpc , the bubble has a size of @xmath35 pc and a hi  mass between 700 and 2400 m@xmath6 . by simple estimates of energetics , ages and stellar wind parameters , we think that this feature is a wind - blown bubble produced by the o - progenitor of the wr star . the presence of a nearly spatially - coincident iras shell reinforces our hypothesis . we have also mapped the co j = @xmath0  and @xmath1  emission adjacent to ngc2359  and found three co components in this region . two of these components , with radial velocities of 34 ( component a1 ) and 67 ( component a2 ) km s@xmath3 , have narrow profiles ( up to 2 km s@xmath3 ) and do not show any morphological or kinematical effects which indicate the disturbance by the nebula or the wr star . however , the third component ( the component b , emitting at 54 km s@xmath3 ) , clearly bounds ngc2359  at its southern and eastern border . the profiles are broad , with linewidths of 45.5 km s@xmath3 . a velocity gradient of a few km s@xmath3  is noted towards the south and south - eastern interface with the nebula . a weak `` bridge '' in velocity of small angular extension seems to connect this component with the component b2 at the southern part . we have observed the @xmath4co j = @xmath0  line at a few points in the region where the component b is detected . we have found a striking variation of the @xmath15co/@xmath4co ratio both in position and in velocity . this ratio seems to decrease near the peak of the co emission , for velocities higher than 53 km s@xmath3 . this particular emiting region is not resolved by the beam , and is probably 0.2 pc wide .    from the observational data , we have estimated the physical conditions for each co component using the lvg approximation . the highest excitation temperature and density are found in the component b. it is remarkable that high temperatures ( up to 80k ) are derived for the co emission projected onto the hii  region .    in view of the kinematics , the morphology and the physical properties of the co , we think that the component a2 was shocked and accelerated by the expanding bubble to reach the radial velocities of the component b. the shock front is still acting at the southern part of the region . the origin of such shock might be related to the wr phase of hd56925  or , more probably , to previous episodes of rsg of lbv . from the @xmath15co/@xmath4co ratio we infer the presence of a thin layer of dense gas which represent the more recently shocked material . our data , however , can not resolve spatially the region where a thin layer of shocked material is expected . co observations with better resolution and sensitivity is needed to disclose the main features of the region . a1 & 0.5 &  & 5 & 0.6 & 1.0 & 0.6 & 0.6 &  & 96 + a2 & 0.5 &  & 6 & 1.0 & 0.4 & 0.8 & 0.8 &  & 28 + b@xmath22 & 0.8 &  & 8 & 1.8 & 0.5 & 0.3 & 0.4 &  & 160 + b@xmath24 & 1.0 & 30 & 10 & 5.0 & 1.6 & 0.6 & 0.8 & @xmath45 & 10 + b@xmath25 & 1.0 & 5 & 10 & 58 & 21 . & 12 &  & @xmath13 0.1 & 120 + b@xmath23 & 2.0 & @xmath17 30 & @xmath17 80 & @xmath46 & 0.1 & @xmath280.1 & 0.2 & @xmath47 0.1 & @xmath48 8 +     + the lvg approximation was used in all cases but b@xmath25 . for this particular region , we have + estimated the size of the emiting region and taken into account the beam filling factor . + all masses were computed by assuming a distance of 5 kpc .
ngc2359  is an optical nebula excited by the powerful wind and the radiation of the wolf 
 rayet star hd56925 . 
 we have investigated the interaction between this massive star and the surrounding neutral gas by analyzing the large - scale 21cm - hi  emission and by mapping the nebula in the j = @xmath0  and the j = @xmath1  lines of co. we found a conspicuous ( @xmath2 pc ) hi  shell , expanding at 12 km s@xmath3 , likely produced during the main - sequence phase of the star . 
 the molecular gas towards ngc2359  shows three velocity components . 
 two of these components , a1 and a2 , have narrow linewidths ( 12 km s@xmath3 ) and radial velocities of 3538 and 6468 km s@xmath3 , respectively . 
 the third component is detected at radial velocities between 50 and 58 km s@xmath3  and has a broader profile ( up to 5.5 km s@xmath3 ) . 
 furthermore , this component is morphologicaly related with the nebula and has a velocity gradient of a few km s@xmath3 . we have also estimated the physical parameters of the molecular gas by means of a lvg modelling of the co emission . 
 the gas projected onto the southern hiiregion of the nebula has low co column density and is rather hot , probably up to 80k . 
 several profiles of the @xmath4co j = @xmath0  line near the peak of the emission , together with a weak emission bridge between the broad and one of the narrow components ( component a2 ) , suggest the presence of a shock front acting in the southern part of the nebula . 
 this shock was likely produced in a previous rsg / lbv phase of hd56925 .